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483060	[ "Grillo, C.", "Gobat, R.", "Presotto, V.", "Balestra, I.", "Mercurio, A.", "Rosati, P.", "Nonino, M.", "Vanzella, E.", "Christensen, L.", "Graves, G.", "Biviano, A.", "Lemze, D.", "Bartelmann, M.", "Benitez, N.", "Bouwens, R.", "Bradley, L.", "Broadhurst, T.", "Coe, D.", "Donahue, M.", "Ford, H.", "Infante, L.", "Jouvel, S.", "Kelson, D.", "Koekemoer, A.", "Lahav, O.", "Medezinski, E.", "Melchior, P.", "Meneghetti, M.", "Merten, J.", "Molino, A.", "Monna, A.", "Moustakas, J.", "Moustakas, L. A.", "Postman, M.", "Seitz, S.", "Umetsu, K.", "Zheng, W.", "Zitrin, A." ]	2014ApJ...786...11G	[ "CLASH: Extending Galaxy Strong Lensing to Small Physical Scales with Distant Sources Highly Magnified by Galaxy Cluster Members" ]	13	[ "Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 Copenhagen, Denmark", "Laboratoire AIM-Paris-Saclay, CEA/DSM-CNRS-Universitè Paris Diderot, Irfu/Service d'Astrophysique, CEA Saclay, Orme des Merisiers, F-91191 Gif sur Yvette, France", "Dipartimento di Fisica, Università degli Studi di Trieste, via G. B. Tiepolo 12, I-34143 Trieste, Italy", "INAF-Osservatorio Astronomico di Trieste, via G. B. Tiepolo 11, I-34131 Trieste, Italy; INAF-Osservatorio Astronomico di Capodimonte, Via Moiariello 16, I-80131 Napoli, Italy", "INAF-Osservatorio Astronomico di Capodimonte, Via Moiariello 16, I-80131 Napoli, Italy", "Dipartimento di Fisica e Scienze della Terra, Università degli Studi di Ferrara, Via Saragat 1, I-44122 Ferrara, Italy", "INAF-Osservatorio Astronomico di Trieste, via G. B. Tiepolo 11, I-34131 Trieste, Italy", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 Copenhagen, Denmark", "Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA", "INAF-Osservatorio Astronomico di Trieste, via G. B. Tiepolo 11, I-34131 Trieste, Italy", "Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA", "Institut für Theoretische Astrophysik, Zentrum für Astronomie, Universität Heidelberg, Philosophenweg 12, D-69120 Heidelberg, Germany", "Instituto de Astrofisica de Andalucia (CSIC), Glorieta de la Astronomia s/n, E-18008 Granada, Spain", "Leiden Observatory, Leiden University, P.O. Box 9513, NL-2333 Leiden, The Netherlands", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21208, USA", "Department of Theoretical Physics, University of the Basque Country, P.O. Box 644, E-48080 Bilbao, Spain", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21208, USA", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA", "Departamento de Astronomia y Astrofisica, Pontificia Universidad Catolica de Chile, V. Mackenna 4860, Santiago 22, Chile", "Institut de Cincies de l'Espai (IEE-CSIC), E-08193 Bellaterra (Barcelona), Spain; Department of Physics and Astronomy, University College London, Gower street, London WC1E 6BT, UK", "Observatories of the Carnegie Institution of Washington, Pasadena, CA 91101, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21208, USA", "Department of Physics and Astronomy, University College London, Gower street, London WC1E 6BT, UK", "Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA", "Center for Cosmology and Astro-Particle Physics, The Ohio State University, 191 West Woodruff Avenue, Columbus, OH 43210, USA", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA", "Instituto de Astrofisica de Andalucia (CSIC), Glorieta de la Astronomia s/n, E-18008 Granada, Spain", "Institut für Astronomie und Astrophysik, Universitäts-Sternwarte München, D-81679 München, Germany", "Department of Physics and Astronomy, Siena College, 515 Loudon Road, Loudonville, NY 12211, USA", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21208, USA", "Institut für Astronomie und Astrophysik, Universitäts-Sternwarte München, D-81679 München, Germany", "Institute of Astronomy and Astrophysics, Academia Sinica, P.O. Box 23-141, Taipei 10617, Taiwan", "Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, MS249-17, Pasadena, CA 91125, USA; Hubble fellow" ]	[ "2015A&A...574A..11K", "2015A&A...579A...4G", "2015ApJ...800...38G", "2016ApJ...820...43S", "2016MNRAS.458.1493P", "2017MNRAS.470...95M", "2019A&A...631A.130B", "2020A&A...635A..98R", "2021A&A...648A.123B", "2022A&A...668A.188M", "2023A&A...679A.124G", "2024A&A...681A..68E", "2024SSRv..220...19N" ]	[ "astronomy" ]	4	[ "dark matter", "galaxies: clusters: individual: MACS J1206.2–0847", "galaxies: high-redshift", "galaxies: stellar content", "galaxies: structure", "gravitational lensing: strong", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]	[ "1955ApJ...121..161S", "1976ApJ...204..668F", "1999MNRAS.310..540A", "2000ApJ...533..682C", "2000ApJ...536..571B", "2001astro.ph..2340K", "2001astro.ph..2341K", "2003ApJ...595...29R", "2003MNRAS.344.1000B", "2003SPIE.4854...81F", "2004ApJ...610...69K", "2004ApJ...611..739T", "2004ApJS..150....1B", "2004PASP..116..138C", "2005ApJ...623..666R", "2006A&A...457..841I", "2006AJ....132..926C", "2006ApJ...638..703B", "2006ApJ...640..662T", "2006ApJ...649..599K", "2006MNRAS.371..703S", "2007A&A...461..881L", "2007ApJ...656..739H", "2007ApJS..172..196K", "2008A&A...477..397G", "2008A&A...477L..25G", "2008A&A...486...45G", "2008ApJ...684.1075M", "2008SPIE.7010E..1EK", "2009A&A...501..461G", "2009A&A...502..445L", "2009ApJ...690..670T", "2009ApJ...703L..51K", "2009ApJ...705.1099A", "2009MNRAS.392..104B", "2009MNRAS.399...21B", "2010A&A...524A..94S", "2010AJ....139.2097P", "2010ARA&A..48...87T", "2010ApJ...708..750S", "2010ApJ...709.1195T", "2010ApJ...710..372G", "2010ApJ...711..201S", "2010ApJ...711..246F", "2010ApJ...721L.163A", "2010ApJ...722..779G", "2010ApJ...724..511A", "2010CQGra..27w3001B", "2010MNRAS.402L..44R", "2010MNRAS.402L..67G", "2010MNRAS.406.1220W", "2010MNRAS.408.1969V", "2011A&A...528A..73D", "2011A&A...532A..95F", "2011ApJ...727...96R", "2011ApJ...742..117Z", "2011ApJS..197...36K", "2011MNRAS.415.2215B", "2011MNRAS.417.3000S", "2011MNRAS.418..929G", "2012ApJ...744...41B", "2012ApJ...749...15G", "2012ApJ...749...38M", "2012ApJ...749...97Z", "2012ApJ...752..163S", "2012ApJ...753L..32S", "2012ApJ...755...56U", "2012ApJ...757...22C", "2012ApJ...757...82B", "2012ApJ...761..170G", "2012ApJS..199...25P", "2012MNRAS.419..936F", "2012MNRAS.427.1918E", "2012Natur.481..341V", "2012Natur.484..485C", "2012Natur.489..406Z", "2013A&A...558A...1B", "2013A&A...559L...9B", "2013ApJ...762...32C", "2013ApJ...766...70S", "2013ApJ...773....7D", "2013ApJ...774..124E", "2013ApJ...776...91L", "2013ApJ...777...43M", "2013ApJ...777...97S", "2013ApJ...777...98S", "2013MNRAS.429L..35A", "2013MNRAS.433.2604G", "2013MNRAS.436..253B", "2013MNRAS.436.1040S", "2013arXiv1307.4220X", "2014A&A...562A..86J", "2014ApJ...783...28G", "2014ApJ...786....9P", "2014ApJ...792...76B", "2014ApJ...795..126B", "2014MNRAS.438.1417M" ]	[ "10.1088/0004-637X/786/1/11", "10.48550/arXiv.1403.0573" ]	1403	1403.0573_arXiv.txt	Gravitational lensing studies have radically improved our understanding of the internal structure of galaxies and clusters of galaxies (e.g., \citealt{tre10b}; \citealt{bar10}). In particular, the combination of strong lensing with stellar dynamics or stellar population synthesis models has allowed to characterize some properties, previously almost unexplored, of the galaxy dark-matter haloes and sub-haloes. For example, it has become possible to measure the dark over total mass fraction and dark-matter halo density slope in the inner regions of galaxies (e.g., \citealt{gri09,gri10c,gri12}; \citealt{aug09}; \citealt{bar09,bar11}; \citealt{son12}; \citealt{eic12}), and to estimate the mass function of dark satellites (also called substructure) (e.g., \citealt{koc04}; \citealt{veg10,veg12}; \citealt{fad12}; \citealt{xu13}) and the spatial extent (e.g., \citealt{hal07}; \citealt{suy10}; \citealt{ric10}; \citealt{don11}; \citealt{eic13}) of galaxy dark-matter haloes. Moreover, the same combinations of mass diagnostics have enabled to investigate the total mass density profile (e.g., \citealt{rus03,rus05}; \citealt{koo06,koo09}; \citealt{ruf11}; \citealt{son13b}; \citealt{agn13}), the stellar Initial Mass Function (IMF; \citealt{gri08a,gri09}; \citealt{tre10}; \citealt{spi11,spi12}; \citealt{bar13}), and the origin of the tilt of the Fundamental Plane (e.g., \citealt{gri09,gri10b}; \citealt{aug10b}) of massive early-type galaxies and to use single lenses or statistical samples of them to infer cosmologically relevant quantities (e.g., \citealt{gri08b}; \citealt{sch10}; \citealt{suy10b,suy13}; \citealt{fad10}). Taking advantage of the excellent data collected by the Lenses Structure and Dynamics (LSD; \citealt{tre04}), Sloan Lens ACS Survey (SLACS; \citealt{bol06}; \citealt{tre06}; \citealt{aug10b}), CFHT Strong Lensing in the Legacy Survey (SL2S; \citealt{mor12}; \citealt{gav12}; \citealt{son13a}), and the BOSS Emission-Line Lens Survey (BELLS; \citealt{bro12}; \citealt{bol12}), the analyses conducted so far have mainly examined the physical properties of isolated, massive early-type galaxies, acting as strong lenses on background sources. Only more recently, thanks also to the Cambridge And Sloan Survey Of Wide ARcs in the skY (CASSOWARY; \citealt{bel09}; \citealt{sta13}), growing interest has been shown in the study of early-type lens galaxies residing in galaxy groups and clusters (e.g., \citealt{gri08c,gri11,gri13}; \citealt{lim09}; \citealt{dea13}). Despite the large amount of results published to date, detailed strong lensing studies in ``small'' lens galaxies are still lacking, mostly because these systems are not observed frequently. One possibility to find more such systems is to look at clusters of galaxies, where, owing to the increase in the strong lensing cross section of a cluster member due to the presence of the extended mass distribution of the cluster, low-mass galaxies are more likely to produce strong lensing features than in less dense environments. Investigations of these objects are particularly useful, as they can provide the necessary piece of information to elucidate what is the amount and distribution of dark matter in astrophysical objects extending from the lowest to the highest ends of the galaxy mass function. The comparison of these observational measurements over a wide range of physical scales with the outcomes of cosmological simulations can give fundamental clues about the precise nature of dark matter and the role played by the interaction of baryons and dark matter during the mass assembly of cosmological structures. The Cluster Lensing And Supernova survey with Hubble (CLASH; GO 12065, PI Postman) was awarded 524 orbits of Hubble Space Telescope (HST) time to observe 25 massive (virial mass $M_{\mathrm{vir}} \approx 5$-$30 \times 10^{14} M_{\odot}$, X-ray temperature $T_{X} \ge 5$ keV) galaxy clusters in 16 broadband filters, ranging from approximately 2000 to 17000 \AA $\,$ with the Wide Field Camera 3 (WFC3; \citealt{kim08}) and the Advanced Camera for Surveys (ACS; \citealt{for03}). The sample, spanning a wide redshift range ($z$ = 0.18-0.90), was carefully chosen to be largely free of lensing bias and representative of relaxed clusters, on the basis of their symmetric and smooth X-ray emission profiles (for a thorough overview, see \citealt{pos12}). CLASH has four main scientific goals: 1) measure the cluster total mass profiles over a wide radial range, by means of strong and weak lensing analyses (e.g., \citealt{zit11}; \citealt{coe12}; \citealt{med13}); 2) detect new Type Ia supernovae out to redshift z $\sim$ 2.5, to improve the constraints on the dark energy equation of state (e.g., \citealt{gra13}; \citealt{pat13}); 3) discover and study some of the first galaxies that formed after the Big Bang (z $>$ 7) (e.g., \citealt{zhe12}; \citealt{coe13}; \citealt{bow13}); 4) perform galaxy evolution analyses on cluster members and background galaxies. Ancillary science that can surely be carried out with the superb data set of CLASH is the analysis of several new strong lensing systems on galaxy scale. A Large Programme (186.A-0798, PI Rosati) of 225 hours with the VIMOS instrument at the Very Large Telescope (VLT) has also been approved to perform a panoramic spectroscopic survey of the 14 CLASH clusters that are visible from ESO-Paranal (Rosati et al. 2014, in preparation). This observational campaign aims at measuring in each cluster the redshifts of 1) approximately 500 cluster members within a radius of more than 3 Mpc; 2) 10-30 lensed multiple images inside the HST field of view, including possible highly-magnified candidates out to z $\approx$ 7 (e.g., \citealt{mon13}; \citealt{bal13}); 3) possible supernova hosts. In one of the CLASH clusters (i.e., MACS J1206.2$-$0847, hereafter MACS 1206), the first spectroscopic redshifts have already been exploited to build robust strong lensing models (\citealt{zit12}; \citealt{ume12}), to obtain an independent total mass estimate from the spatial distribution and kinematics of the cluster members (\citealt{biv13}; \citealt{lem13}), and to confirm a source at $z = 5.703$ (\citealt{bra13}). Strong lensing (with spectroscopically confirmed systems) and cluster dynamics analyses are planned for all 14 southern clusters (e.g., in MACS J0416.1$-$2403, Grillo et al. 2014, in preparation; Balestra et al. 2014, in preparation). Here, we focus on a rare strong lensing system in which two angularly close early-type galaxies, members of the galaxy cluster MACS 1206 at $z=0.44$, produce in total ten multiple images of a double source located at $z \approx 3.7$. This is the first example of the kind of strong lensing studies that can be conducted on galaxy cluster members, capitalizing on the extraordinary multi-band photometric observations obtained as part of the CLASH program and spectroscopic measurements of the VLT/VIMOS follow-up campaign. This work is organized as follows. In Sect. 2, we introduce the photometric and spectroscopic observations used for this analysis. In Sect. 3, we present the strong lensing modeling performed to measure principally the total mass values of the lens galaxies and the magnification factors of the multiple images. In Sect. 4, we estimate the luminous mass values of the lens and lensed galaxies by means of stellar population synthesis models. In Sect. 5, we compare some physical quantities, related to the luminous and total masses, of the two lens galaxies with those of lower and higher-mass galaxies. In Sect. 6, we summarize our conclusions. All quoted errors are 68.3\% confidence limits (CL) unless otherwise stated. Throughout this work we assume $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{m}=0.3$, and $\Omega_{\Lambda}=0.7$. In this model, 1\arcsec$\,$ corresponds to a linear size of 5.69 kpc at the cluster redshift of $z=0.44$.	The combination of unprecedented HST multi-wavelength observations and VLT spectra has allowed us to perform a detailed strong lensing and stellar population analysis of an unusual system composed in total of ten multiple images of a double source, lensed by two early-type galaxies in the field of the CLASH galaxy cluster MACS 1206. Our main results can be summarized in the following points. \begin{itemize} \item[$\bullet$] Based on our 16-band photometry and low-resolution spectroscopy, we measure a photometric redshift of 3.7 for the source and spectroscopic redshifts of 0.436 and 0.439 for the two lens galaxies G1 and G2, respectively, thus confirming their membership to MACS 1206. \item[$\bullet$] By modeling the total mass distribution of the cluster members and cluster in terms of singular isothermal profiles, we can reconstruct well the observed positions of the multiple images and predict a total magnification factor of approximately 50 for the source. \item[$\bullet$] From the lensing modeling statistics, we estimate effective velocity dispersion values of $97 \pm 3$ and $240 \pm 6$ km s$^{-1}$, corresponding to total mass values projected within the effective radii of $1.7_{-0.1}^{+0.1}\times10^{10}$ and $2.8_{-0.1}^{+0.2}\times10^{11}$~$M_{\odot}$ for G1 and G2, respectively. Moreover, we obtain reasonable values for the distribution and amount of projected total mass in the galaxy cluster component. \item[$\bullet$] Through composite stellar populations synthesis models (adopting a Salpeter stellar IMF), we infer luminous mass values of $(1.7 \pm 0.7)\times10^{10}$ and $(4.5 \pm 1.8)\times10^{11}$ $M_{\odot}$ for, respectively, G1 and G2, and $(1.0 \pm 0.5)\times10^{9}$ $M_{\odot}$ for the source, taking into account the estimated lensing magnification factor. \item[$\bullet$] In G1 and G2, respectively, we derive luminous over total mass fractions of $0.51 \pm 0.21$ and $0.80 \pm 0.32$. We compare these values with those typical of massive early-type galaxies and dwarf spheroidals and conclude that more analyses in the CLASH fields of systems similar to that presented here will enable us to extend the investigation of the internal structure of galaxies in an important and still relatively unexplored region of the physical parameter space. \end{itemize}	14	3	1403.0573	We present a complex strong lensing system in which a double source is imaged five times by two early-type galaxies. We take advantage in this target of the extraordinary multi-band photometric data set obtained as part of the Cluster Lensing And Supernova survey with Hubble (CLASH) program, complemented by the spectroscopic measurements of the VLT/VIMOS and FORS2 follow-up campaign. We use a photometric redshift value of 3.7 for the source and confirm spectroscopically the membership of the two lenses to the galaxy cluster MACS J1206.2-0847 at redshift 0.44. We exploit the excellent angular resolution of the HST/ACS images to model the two lenses in terms of singular isothermal sphere profiles and derive robust effective velocity dispersion values of 97 ± 3 and 240 ± 6 km s<SUP>-1</SUP>. Interestingly, the total mass distribution of the cluster is also well characterized by using only the local information contained in this lensing system, which is located at a projected distance of more than 300 kpc from the cluster luminosity center. According to our best-fitting lensing and composite stellar population models, the source is magnified by a total factor of 50 and has a luminous mass of approximately (1.0 ± 0.5) × 10<SUP>9</SUP> M <SUB>⊙</SUB> (assuming a Salpeter stellar initial mass function). By combining the total and luminous mass estimates of the two lenses, we measure luminous over total mass fractions projected within the effective radii of 0.51 ± 0.21 and 0.80 ± 0.32. Remarkably, with these lenses we can extend the analysis of the mass properties of lens early-type galaxies by factors that are approximately two and three times smaller than previously done with regard to, respectively, velocity dispersion and luminous mass. The comparison of the total and luminous quantities of our lenses with those of astrophysical objects with different physical scales, like massive early-type galaxies and dwarf spheroidals, reveals the potential of studies of this kind for improving our knowledge about the internal structure of galaxies. These studies, made possible thanks to the CLASH survey, will allow us to go beyond the current limits posed by the available lens samples in the field. <P />This work is based on data collected at NASA HST and at ESO VLT (Program ID 186.A-0798 and 089.A-0879).	false	[ "galaxies", "robust effective velocity dispersion values", "total mass fractions", "lens early-type galaxies", "massive early-type galaxies", "dwarf spheroidals", "a Salpeter stellar initial mass function", "redshift", "s", "Supernova survey", "singular isothermal sphere profiles", "ESO VLT", "the galaxy cluster", "different physical scales", "factors", "a luminous mass", "186.A-0798", "two early-type galaxies", "studies", "NASA HST" ]	12.872774	4.568682	163
745882	[ "Rodríguez Castillo, Guillermo A.", "Israel, Gian Luca", "Esposito, Paolo", "Pons, José A.", "Rea, Nanda", "Turolla, Roberto", "Viganò, Daniele", "Zane, Silvia" ]	2014MNRAS.441.1305R	[ "Pulse phase-coherent timing and spectroscopy of CXOU J164710.2-45521 outbursts" ]	18	[ "Dipartimento di Fisica, Sapienza Università di Roma, Piazzale Aldo Moro 5, I-00185 Rome, Italy; INAF - Astronomical Observatory of Rome, via Frascati 33, I-00040 Monte Porzio Catone, Italy", "INAF - Astronomical Observatory of Rome, via Frascati 33, I-00040 Monte Porzio Catone, Italy", "INAF - IASF Milano, Via E. Bassini 15, I-20133 Milano, Italy", "Departament de Física Aplicada, Universitat d'Alacant, Ap. Correus 99, E-03080 Alacant, Spain", "Institute of Space Sciences (CSIC-IEEC), Campus UAB, Faculty of Science, Torre C5-parell, E-08193 Barcelona, Spain; Astronomical Institute `Anton Pannekoek', University of Amsterdam, Postbus 94249, NL-1090 GE Amsterdam, the Netherlands", "Department of Physics, University of Padova, Via Marzolo 8, I-35131 Padova, Italy; Mullard Space Science Laboratory, University College London, Holmbury St. Mary, Dorking, Surrey RH5 6NT, UK", "Departament de Física Aplicada, Universitat d'Alacant, Ap. Correus 99, E-03080 Alacant, Spain; Institute of Space Sciences (CSIC-IEEC), Campus UAB, Faculty of Science, Torre C5-parell, E-08193 Barcelona, Spain", "Mullard Space Science Laboratory, University College London, Holmbury St. Mary, Dorking, Surrey RH5 6NT, UK" ]	[ "2015ApJ...805...81W", "2015RPPh...78k6901T", "2016AIPC.1701h0006K", "2016ApJ...828L..13R", "2016MNRAS.456.4145R", "2016MNRAS.458.2088P", "2018Ap&SS.363..184K", "2018MNRAS.473.3180T", "2018MNRAS.474..961C", "2019ApJ...877L..10A", "2019MNRAS.484.2931B", "2019MNRAS.485....2C", "2019MNRAS.485.4274H", "2019MNRAS.487.1426B", "2020ApJ...904L..21Y", "2021ASSL..461...97E", "2021ApJ...920L...4E", "2023JKAS...56...41S" ]	[ "astronomy" ]	5	[ "stars: individual: CXOU J164710.2-455216", "stars: magnetars", "stars: neutron", "X-rays: bursts", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1990ApJ...365L...9P", "1992ApJ...392L...9D", "1995A&A...299L..41V", "1995MNRAS.275..255T", "2000ApJ...535L..55W", "2002ApJ...574..332T", "2003ApJ...599..485D", "2004ApJ...603L..97I", "2004ApJ...609L..21I", "2005MPLA...20.2799H", "2006ATel..893....1C", "2006ApJ...636L..41M", "2006ApJ...648L..51S", "2006MNRAS.372.1407C", "2006Natur.442..892C", "2007AdSpR..40.1453X", "2007ApJ...664..448I", "2007ApJ...666L..93C", "2007MNRAS.378L..44M", "2008A&ARv..15..225M", "2008ApJ...686.1245R", "2008MNRAS.386.1527N", "2009A&A...498..195B", "2009ATel.1909....1I", "2009ApJ...703.1044B", "2010A&A...516A..78N", "2010A&A...516A..88O", "2010ApJ...722..788A", "2011ASSP...21..247R", "2011ATel.3653....1I", "2011ApJ...726...37W", "2011ApJ...727L..51P", "2011ApJ...740..105T", "2012ApJ...748L..12R", "2012ApJ...750L...6P", "2012ApJ...754...27R", "2012PASJ...64...56M", "2013ApJ...763...82A", "2013ApJ...768..144T", "2013ApJ...770...65R", "2013BrJPh..43..356M", "2013IAUS..290...93A", "2013IJMPD..2230024T", "2013MNRAS.434..123V", "2013Natur.500..312T" ]	[ "10.1093/mnras/stu603", "10.48550/arXiv.1403.7165" ]	1403	1403.7165_arXiv.txt	Soft $\gamma$-repeaters (SGRs) and anomalous X-ray pulsars (AXPs) are isolated neutron stars (INSs) with prominent high-energy manifestations. They are characterized by rotational periods in the 0.3--12\,s range and period derivatives (usually) larger than those typical of the radio-pulsar population ($\dot{P}\sim10^{-13}-10^{-10}$s/s). They exhibit peculiar flaring activity (see e.g. Mereghetti 2013) over a large range of time-scales (milliseconds to minutes) and luminosities ($L\sim10^{38-47}$ erg s$^{-1}$). Estimates of their magnetic field, derived under the usual assumptions for isolated rotation-powered pulsars, place them at the high end of the pulsar population ($B\approx 10^{14-15}$\,G). This, and other direct (Tiengo et al. 2013) and indirect evidences, suggests that they host an ultramagnetized neutron star, or magnetar (Duncan \& Thompson 1992, Thompson \& Duncan 1995). Since the detection of SGRs/AXPs as persistent X-ray sources, one of the main concern has been the imbalance between the emitted luminosity and the rotational energy loss rate, $\dot E$. Rotation is believed to be the standard mechanism that provides the energy output in canonical radio-pulsars. However, in SGRs/AXPs $\dot E$ is orders of magnitude below $L_{\mathrm{X}}$, although in some transient sources the rotational energy loss rate may exceed luminosity in the quiescent state (see e.g. Rea et al. 2012a). Energy might be supplied by accretion, if a feeding companion is present as is the case of many X-ray (binary) pulsars. Despite intensive searches, however, no binary companions have been detected so far around SGRs/AXPs (see e.g. Woods et al. 2000; Camilo et al. 2006 for the tightest constraints). A more likely alternative is that SGRs/AXPs are magnetically powered sources in which the decay/rearrangement of their (huge) magnetic field is responsible for both their persistent and bursting emission. Nowadays the magnetar model appears to be the more viable one and it will be assumed in this investigation, in particular for what concerns the timing and spectral analysis. In the following we shall refer to SGRs/AXPs as the ``magnetar candidates'', or simply as magnetars. Alternative scenarios have been proposed with varying degree of success to explain the SGRs/AXPs phenomenology, and include fallback discs (see e.g. Alpar et al. 2012), Thorne-{\.Z}ytkow objects (van Paradijs, Taam, \& van den Heuvel 1995), strange/quark/hybrid stars (see e.g. Horvath 2005; Xu 2007; Ouyed, Leahy \& Niebergal 2010) and fast rotating, highly-magnetized (B $\sim 10^{8-9}$ G) massive white dwarfs (Paczy\'nski 1990; Malheiro, Rueda \& Ruffini 2012), among others (see Turolla \& Esposito 2013, section 5 and Mereghetti 2008, section 7 for overviews). \subsection{OUTBURSTS IN MAGNETARS} Most of the known magnetar candidates are transient sources. A transient episode in a magnetar can be defined as an outburst characterized by a rapid (minutes--hours) increase of the persistent flux by a factor of $\sim 10$--1000, with a subsequent decay back to the quiescent level on time-scales of months--years. Short bursts, which usually trigger detection, are emitted in the early phases of the outburst. Recurrent outbursts have been observed in a few sources (see Rea \& Esposito, 2011 for a review). Within the magnetar picture, outbursts occur quite naturally. According to our current understanding, one of the major differences between the magnetar candidates and pulsars is not (or not only) the higher value of the dipole field (there are low-field magnetars with $B\la 10^{13}$ G and high-field pulsars with $B\ga 10^{13}$ G), but the presence of a strong toroidal component in the internal field (Turolla et al. 2011 and references therein). It is the dissipation of the internal field which powers the magnetar bursting/oubursting behaviour by injecting energy deep in the star crust and/or by inducing displacements of the surface layers, with the consequent ``twisting'' of the external field (e.g. Thompson, Lyutikov \& Kulkarni 2002; Perna \& Pons 2011; Pons \& Rea 2012). The rate at which these episodes occur is different in different sources and is believed to depend mostly on the star magnetic field at birth and on its age. Since the discovery the first confirmed transient magnetar in 2003 (XTE J1810-197, which exhibited a flux enhancement by a factor of $> 100$; Ibrahim et al. 2004), outbursts have been the object of much interest. This stems from the possibility of testing, during the outburst decay, theoretical predictions over a relatively large luminosity range in a single source, where a large number of important parameters are not changing, like, e.g., distance, mass, radius, age, viewing geometry (see e.g. Bernardini et al. 2009; Albano et al. 2010; Rea et al 2013). \src\ is among the transients with the larger flux variation. Following the outburst of 2006 September, in fact, its flux grew by a factor of $\gtrsim 300$ (Campana \& Israel 2006). \begin{table} \renewcommand{\arraystretch}{1.3} \resizebox{8cm}{!} { \begin{minipage}{7.5cm} \begin{tabular}{\|l\|l\|c\|c\|c\|} \hline \renewcommand{\arraystretch}{1.0} \multirow{2}{}{Telescope} & \multirow{2}{}{Date\footnote{Start of observation (post-reduction).} (MJD TDB)} & Exposure & \multirow{2}{}{Observation ID} & Name \\ & & time (ks)& & ([t]YYMMDD)\\ \hline \renewcommand{\arraystretch}{1.5} \multirow{2}{}{\XMM\footnote{In all \XMM\ observations EPN detector was used.}} &53994.791448810 & 46.0 & 0404340101 & 060916\\ & 54000.527619667 & 29.2 & 0311792001&060922\\\hline \multirow{5}{}{\emph{Chandra}\footnote{In all \emph{Chandra} data ACIS detector was used.}} & 54005.283617719 & 15.7 & 6724 & c060927\\ & 54009.985545836 & 20.7 & 6725 & c061002\\ & 54017.265934068 & 26.2 & 6726 & c061009 \\ & 54036.293110632 & 15.7 & 8455 & c061028\\ & 54133.801451742 & 20.6 & 8506 & c070113\\ \hline \multirow{5}{}{\XMM$^b$} & 54148.378135941 & 17.6 & 0410580601& 070217\\ & 54331.412027612 & 23.7 & 0505290201& 070819\\ & 54511.305592759 & 29.8 & 0505290301& 080215\\ & 54698.508875675 & 30.7 & 0555350101& 080820\\ & 55067.333418577 & 41.4 & 0604380101& 090824\\ \hline \multirow{2}{}{\emph{Swift/PC}\footnote{In all \emph{Swift} observations we refer to the XRT}} & 55823.887347023 & 3.1 & 00030806020 & \\ & 55829.233491897 & 4.3 & 00030806022 & \\ \hline \XMM$^b$ & 55831.936336540 & 16.7& 0679380501& 110927\\ \hline \multirow{5}{}{\emph{Swift/PC}$^d$} & 55835.185790895 & 3.7 & 00030806023 & \\ & 55839.066673333 & 3.7 & 00030806024 & \\ & 55842.093900977 & 3.9 & 00030806025 &\\ & 55844.092375078 & 4.0 & 00030806026 & \\ & 55849.040224474 & 8.8 & 00030806027 & \\ \hline \emph{Chandra}$^c$ & 55857.646333201 & 19.1 & 14360 & c111023 \\ \hline \multirow{2}{}{\emph{Swift/PC}$^d$} & 55974.223899273 & 0.6 & 00030806028-29 \\ & 56001.015943907 & 2.4 & 00030806031 \\ \hline \hline \end{tabular} \end{minipage} } \caption{Summary of the observational data used in the paper} \label{data} \end{table*} \begin{figure} \centering \includegraphics[width=.34\textwidth, angle=270]{fig1.ps} \caption{Upper panel: pulse profile of \src, observation 060922 at the $0.5 - 2.0$ keV energy range. Lower panel: The same observation at the $3.0 - 12.0$ keV energy range. Note that the peak at phase $\sim 0.4$ is missing at higher energies.} \label{ene} \end{figure} \subsection{CXOU J164710.2-455216} The source, CXOU J1647-45 for short, was discovered by Muno et al. (2006) using \chandra\ observations, with a period of 10.6107(1) s. An important feature of \cxou\ is that it very likely belongs to the young, massive Galactic starburst cluster Westerlund 1. This provides hints about its progenitor and also about its distance. Indeed, studies of the massive stellar population of Westerlund 1 indicate a distance of $\sim 5.0$ kpc and a progenitor with an initial mass $M_i>40\:M_{\odot}$ (Crowther et al. 2006; Muno et al. 2006 Negueruela Clark \& Ritchie 2010). Another prominent feature of \cxou, as mentioned before, is that it underwent an outburst with one of the largest flux enhancement observed up to now among the magnetars. On September 2006 the Burst Alert Telescope (BAT) on board the \swift\ satellite detected an intense burst in the direction of the Westerlund 1. A second observation, performed 13 h later by \swift, with its narrow field instrument, the X-ray telescope (XRT), found \cxou\ brighter by a factor of $\sim 300$. Between February 2007 and August 2009 we requested and obtained five \emph{XMM-Newton} pointings which, together with the September 2006 post-outburst observations, were aimed at studying the evolution of the timing and spectral properties of \cxou\ over a range covering a factor of about 50 in flux, from $\sim 10^{35}$ \ergs\ down to near the quiescent level, at a few $10^{33}$ \ergs\ (Campana \& Israel 2006). Deep observational campaign in the radio, near-infrared and hard X-ray bands did not detect any convincing counterpart (Muno et al. 2006; Israel et al. 2007), in contrast with the results obtained for other transient magnetars, e.g. XTE J1810 (Israel et al. 2004; Camilo et al. 2006) and 1E 1547 (Camilo et al. 2007; Israel et al. 2009). On 2011 September 19 \swift-BAT recorded four relatively bright bursts from a position consistent with that of \cxou\ (Baumgartner et al. 2011), approximately five years after the 2006 outburst onset. A subsequent \swift-XRT pointing found \cxou\ at a flux level of $\sim 7.8 \times 10^{-11}$ \ergscms, more than 200 times higher than its quiescent level ($2.7 \times 10^{-13}$ \ergscms, Muno et al. 2007), and more than 100 times brighter than the latest \emph{XMM-Newton} pointing of August 2009 (Israel, Esposito \& Rea 2011): the pulsar entered a new outburst phase. Several \swift\ observations were requested together with two director's discretionary time observations, one with \emph{XMM-Newton} and one with \emph{Chandra}. The latter two were carried out 9 and 34 days after the BAT trigger, respectively. A further \emph{XMM-Newton} pointing scheduled for April 2012 was cancelled because of a strong solar storm. The \emph{XMM-Newton} and \emph{Chandra} pointings aimed at comparing the properties of the 2006 and 2011 outbursts. \cxou\ 2006 outburst has been analyzed in previous investigations. In particular timing and spectral analyses have been performed by Israel et al. (2007) and Woods et al. (2011; both phase-coherent), and An et al. (2013; period evolution). Their timing results are summarized in Table \ref{altime}. The phase-averaged fluxes and periods for the 2011 outburst, as derived from the \emph{XMM-Newton} and \emph{Chandra} pointings, were reported by An et al. (2013). A detailed spectral and timing analysis is first reported in this paper where we present an extended, phase coherent long-term timing solution and phase-resolved spectroscopic analysis for both outbursts. The implications, within the magnetar scenario, are also discussed by means of state-of-the-art magnetothermal evolution simulations.	\label{discussion} \subsection{Timing} Significant changes in the pulse profile during the outburst decay mean that peaks identification and the way of taking into account their variations in relative phase (among peaks), intensity and shape, is important in order to successfully phase connect the observations. For instance Woods et al. (2011) cite the ``extreme change in pulse profile'' as the reason why they were not able to phase connect the 070819 observation with their coherent solution. On the other hand, An et al. (2013) cite a large time separation between 070819 and the previous observation as the cause of their phase connection loss. As mentioned before (see Fig. \ref{ene}) at different energy ranges the peaks behave differently. This fact, coupled with measurements of the relative phase distances between peaks allowed us to identify them. Once we obtained a consistent peak identification, we had no problems to keep the phase coherence, see Fig. \ref{311}, left panel. We believe that discrepancies with respect to previous published results may be due to the different assumptions used for the phase-fitting algorithm. The new spin-down value $\dot{P}̇ \simeq 9.7 \times 10^{-13}\;$s s$^{-1}$ is similar to that of the two previous $P$ and $\dot{P}$ solutions: $\dot{P} \simeq 9.2 \times 10^{-13}\;$s s$^{-1}$ derived by Israel et al. (2007) and $\dot{P} \simeq 8.3 \times 10^{-13}\;$s s$^{-1}$ reported by Woods et al. (2011) but significantly smaller then the one of the cubic solution of Woods et al. 2011 ($\dot{P} \simeq 13 \times 10^{-13}\;$s s$^{-1}$), who considered a shorter data sample spanning from 2006 September 23 to 2007 February 17. This may be due to a decrease of the spin-down rate throughout the outburst decay. The $P$ and $\dot{P}̇$ values inferred imply a surface dipolar field B $\sim 1.0 \times 10 ^{14}$ G using the conventional formula at the equator B $= 3.2 \times 10 ^{19} (P\dot{P}̇)^{-1/2}$, assuming an orthogonally rotating neutron star of radius 10\,km and moment of inertia 10$^{45}$ g cm$^2$. This estimate lays in the standard magnetar range and agrees with the magnetar nature of \src. \subsection{Outburst decay} In previous (phase-averaged) studies, spectra have been analyzed during the outburst decay and fitted with an absorbed PL plus a BB (Woods et al. 2011; An. et al. 2013). In those works the BB evolution during the outburst agrees with that of the general trend of the warm components of the outburst peaks of our PPS. Particularly, an almost constant temperature % and a shrinking BB radius during the outburst decay (Figs. \ref{pps2} and \ref{pps2b}). To our knowledge, the only other work that have performed a spectral analysis over long time-scales is the one of \citet{Albano2010}, who, in the framework of the twisted magnetospheric model (Thompson et al. 2002), used three modified BB to model the spectra, similar to the approach we based our work on (see text for details). Taking into account that \citet{Albano2010} did not perform a spectral fit, but obtained the physical parameters from synthetic pulse profiles, and, more importantly, that their values correspond to phase-average spectra, it is difficult to make a direct comparison with our results. Nonetheless, the thermal evolution of the BB modeled on \citet{Albano2010} is very similar to the one we infer: a constant warm component and a slightly decreasing value of the hot component temperature, while still consistent with a constant within the $1\sigma$ errors. Likewise, the radius evolution of the hot component in \citet{Albano2010} is as well very similar to the one we infer for the peaks: it significantly decreases throughout the outburst decay, ultimately disappearing in about 700 d. On the other hand, in \citet{Albano2010} the warm component increases in size throughout the outburst, in contrast with what we infer in this work. Yet, the analysis of another magnetar considered by \citet{Albano2010}, XTE J1810-197, show the same decay trend we see in the outburst peaks: the hot and warm components keep an almost constant temperature and fade away in size, leaving the star emitting at the quiescence temperature towards the end of the outburst decay (in the case of a third, constant, ``cool'' BB temperature, see Albano et al. 2010 for details). Prior to the outburst, the pulse profile of the \src\ was single-peaked. The outburst strongly changes the observed pulse profile, and a three-peaked structure is clearly seen from the onset and during most of the outburst. Nevertheless, as the outburst decays, the pulse profile evolves and towards the end of the 2006 campaign, as the luminosity decreases, and \src\ returns to its quiescence level, the pulse profile ``returns'' to a single-peaked structure. The remaining peak correspond to the peak 1, and it is plausibly to assume that it correspond to the quiescence single peak. The radius shrinking decay picture fits into the untwisting magnetosphere (UM) model (Beloborodov, 2009), where current-carrying ``j-bundles'' with twisted magnetic fields gradually shrink. A simple UM model predicts the relation $L \propto A^2$ between the luminosity and the emitting area (see Beloborodov 2009, Equation 48); in Fig. \ref{L.vs.A} we compare the emitting area evolution with luminosity decay for the warm component of the third peak with this modeled relationship. The PL fits well the data but our analysis suggests a somehow flatter dependence then expected by the simple model, see Fig. \ref{L.vs.A}. \begin{figure} \centering \includegraphics[width=0.35\textwidth, angle=270]{fig12.ps} \caption{Luminosity versus emitting area of the warm BB component of the third peak (See Fig. \ref{pps1} and \ref{timing} for reference). The dashed lines represent the $L \propto A^{2}$ of simple untwisting magnetosphere models, see Beloborodov, (2009). The solid line is a PL fit to the data which yields $L \propto A^{1.25}$. The dot-dashed lines represent the $3\sigma$ uncertainty of the fit (and correspond to $L \propto A^{1.17}$ and $L \propto A^{1.34}$).} \label{L.vs.A} \end{figure} An important issue is that this interpretation is model-dependent and modeling the data with other spectral components can potentially yield a different picture. Indeed, other non-purely thermal models may also fit well the data, for instance, a BB+PL model and a resonant cyclotron Scattering (RCS) model (Rea et al. 2008) also fit the data acceptably. For instance, the fit for 060922, the best observation in terms of signal-to-noise ratio, has $\chi^{2}_{red} = 0.97252$ (130 dof) and $\chi^{2}_{red} = 1.0993$ (130 dof), for the BB+PL and the RCS models respectively. While our 1+2BB model has $\chi^{2}_{red} = 1.0210$ (129 dof). On the other hand, independently from the spectral analysis the non-zero, negative second period derivative can also be accounted for within the UM model, as the magnetic field untwists, the spin-down torques diminish, effectively lowering the spin-down rate. However, there may be other explanations to the observed second period derivative, as wind braking, see e.g. \citet{tong2013}. \begin{figure} \includegraphics[width=6.1cm,angle=270]{fig13.ps} \caption{Time evolution of the 0.5-10 keV unabsorbed flux, compared to the predicted light curve of the model discussed in the text. \label{outburstmodel}} \end{figure} Furthermore, we also compared the observations of the 2006 outburst with the theoretical model presented in Pons \& Rea (2012). The pre-outburst model is taken to be the evolved NS that fits the present observational constraints. Then we assume that the source undergoes a sudden starquake, possibly with internal magnetic re-connection, which we model by the injection of energy ($\approx 10^{25}$-$10^{26}$ erg cm$^{-3}$) in the thin layer between two variable densities. We found a good agreement with the luminosity data when the energy is deposited between $\rho = 2 \times 10^{9}$ and $2 \times 10^{10}$ g cm$^{-3}$, precisely in the transition region between the outer crust and the liquid envelope, which may be a hint that the energy is provided not only by elastic energy stored in the solid crust but also by magnetic re-connection in the liquid layer. The time evolution of the unabsorbed flux in the 0.5-10 keV band for this particular model is shown in Fig. \ref{outburstmodel} and superimposed to the measured flux values. The total injected energy was of $2 \times 10^{43}$ erg. We note that the last observation seems to show a smaller flux than the prediction of the theoretical model. Interestingly, the same effect has been observed and discussed for \sgr\ (SGR 0418), where the sudden decrease of the flux after 300 days is not well understood (Rea et al. 2013). The occurrence of a second outburst soon after this last data point, does not allow to determine if the source had fully recovered its quiescence state or it was still cooling down. Note that a short-term ($\sim 10$ d) rise in temperature early in the outburst onset, reported by An et al. (2013), which may be expected from crustal cooling models is out of the long-term evolution scope of this paper. \subsection{Magnetorotational Evolution} \label{magnetorot} \begin{figure} \includegraphics[width=8.3cm,angle=0]{fig14UP.ps} \includegraphics[width=8.3cm,angle=0]{fig14MID.ps} \includegraphics[width=8.3cm,angle=0]{fig14LOW.ps} \caption{From top to bottom, the evolution of the luminosity, period and period derivative according to the model discussed in the text, compared to the measured values. \label{mag-rot-evol}} \end{figure} As previously done for other magnetars (SGR 0418 and Swift J1822.3-1606; see Turolla et al. 2011 and Rea et al. 2012b, 2013) we explore if the magnetothermal evolution of a NS born with standard magnetar conditions can lead to objects with properties compatible with those of CXOU J1647 at the present age. We performed some runs using state-of-the-art magnetothermal evolution codes (see Vigan\`o et al. 2013) assuming a $1.4M_\odot$ NS with radius $R=11.6$ km, a short initial period of 10 ms and initial, purely dipolar field of $B=1.5\times 10^{14}$ G. In the resulting scenario, the Hall term reorganize the internal field, producing a toroidal component of the same strength as the poloidal one on a relatively short time-scale (within a few kyr). We show in Fig. \ref{mag-rot-evol} the evolution of the luminosity, period $P$ and period derivative $\dot P$. The latter two quantities are obtained, from the value of the magnetic field at the equator $B(t)$, by numerical integration of the expression (\citealt{Spitkovsky2006}) \begin{equation} \label{spindown} P \dot{P} \simeq \frac{4B_e^2R^6 \pi^2}{I c^3} (1+\sin^2{\alpha}) \end{equation} where $I$ is the effective moment of inertia of the star, $\alpha$ is the angle between the rotational and the magnetic axis and $c$ is the speed of light. The shaded blue area in the figure includes the uncertainty in the angle $\alpha$. Indeed the properties of CXOU J1647 are recovered by this model at an age between 65 and 90 kyr, about half the spin-down age, which suggests that the magnetic field has not experienced dramatic changes over time. Although the components of the internal initial field $B_{tor}(t=0)$ can be varied to some extent, this would not change our results unless the toroidal field contains most of the magnetic energy ($> 90 \%$), as discussed in Vigan\`o et al. (2013). Moderate values of the initial toroidal field (or higher order poloidal multipoles), are unconstrained and will result in very similar properties at the present age. We can also estimate the current outburst rate of this source following the procedure of Perna \& Pons (2011), which gives $\lesssim 10^{-2}$ yr$^{-1}$. Therefore, within our model, the occurrence of a second outburst in 2009, three years after the first outburst, must be connected to the first event. Since the second outburst is less powerful, the pulse profile after it closely resembles the pulse profile after the initial (2006) one, and the pulsed fraction does not present a strong change (as the sharp fall after the 2006 outburst onset), but rather seems to follow the rising trend seen during the outburst (see Fig. \ref{PF}); it may be speculated that there is a connection between them, of the kind ’main event + sequel’, which could reconcile the model with the observations. \subsection*	14	3	1403.7165	We present a long-term phase-coherent timing analysis and pulse-phase resolved spectroscopy for the two outbursts observed from the transient anomalous X-ray pulsar CXOU J164710.2-455216. For the first outburst we used 11 Chandra and XMM-Newton observations between 2006 September and 2009 August, the longest baseline yet for this source. We obtain a coherent timing solution with P = 10.61065583(4) s, Ṗ = 9.72(1) × 10<SUP>-13</SUP> s s<SUP>-1</SUP> and P̈ = -1.05(5) × 10<SUP>-20</SUP> s s<SUP>-2</SUP>. Under the standard assumptions this implies a surface dipolar magnetic field of ∼10<SUP>14</SUP> G, confirming this source as a standard B magnetar. We also study the evolution of the pulse profile (shape, intensity and pulsed fraction) as a function of time and energy. Using the phase-coherent timing solution we perform a phase-resolved spectroscopy analysis, following the spectral evolution of pulse-phase features, which hints at the physical processes taking place on the star. The results are discussed from the perspective of magnetothermal evolution models and the untwisting magnetosphere model. Finally, we present similar analysis for the second, less intense, 2011 outburst. For the timing analysis we used Swift data together with 2 XMM-Newton and Chandra pointings. The results inferred for both outbursts are compared and briefly discussed in a more general framework.	false	[ "magnetothermal evolution models", "Ṗ = 9.72(1", "=", "-1.05(5", "energy", "the transient anomalous X-ray pulsar CXOU", "a standard B magnetar", "pulsed fraction", "similar analysis", "place", "the untwisting magnetosphere model", "time", "Chandra", "a surface dipolar magnetic field", "Swift data", "pulse-phase features", "intensity", "a coherent timing solution", "10.61065583(4", "a long-term phase-coherent timing analysis and pulse-phase resolved spectroscopy" ]	5.325739	4.494667	69
868276	[ "Liu, Nian-Ping", "Qian, Sheng-Bang", "Liao, Wen-Ping", "He, Jia-Jia", "Zhao, Er-Gang", "Liu, Liang" ]	2014RAA....14.1157L	[ "Photometric investigation of the K-type extremely shallow contact binary V1799 Orion" ]	9	[ "Yunnan Observatories, Chinese Academy of Sciences, Kunming 650011, China; Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650011, China; University of Chinese Academy of Sciences, Beijing 100049, China", "Yunnan Observatories, Chinese Academy of Sciences, Kunming 650011, China; Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650011, China; University of Chinese Academy of Sciences, Beijing 100049, China", "Yunnan Observatories, Chinese Academy of Sciences, Kunming 650011, China; Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650011, China", "Yunnan Observatories, Chinese Academy of Sciences, Kunming 650011, China; Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650011, China", "Yunnan Observatories, Chinese Academy of Sciences, Kunming 650011, China; Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650011, China", "Yunnan Observatories, Chinese Academy of Sciences, Kunming 650011, China; Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650011, China" ]	[ "2016NewA...44...12K", "2019PASJ...71...21K", "2019PASJ...71...81S", "2020MNRAS.497.3493Z", "2020PASJ...72...73L", "2020RAA....20...62A", "2021PASJ...73.1470S", "2021RAA....21..122L", "2023RAA....23b5017P" ]	[ "astronomy" ]	5	[ "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1940BHarO.913....9H", "1951MNRAS.111..642O", "1951PRCO....2...85O", "1967ZA.....65...89L", "1968AJ.....73...26B", "1969AcA....19..245R", "1970VA.....12..217B", "1971ApJ...166..605W", "1974AcA....24..119R", "1976ApJ...205..208L", "1976ApJ...205..217F", "1977MNRAS.179..359R", "1979ApJ...234.1054W", "1981A&A...102...81R", "1985ApJS...58..413B", "1990ApJ...356..613W", "1993AJ....106.2096V", "1994PASP..106..921W", "2000AJ....119.1901A", "2000asqu.book....1C", "2009IBVS.5871....1D", "2009RAA.....9..349Z", "2010AJ....140..215Z", "2010IBVS.5920....1D", "2010IBVS.5929....1N", "2010Obs...130..364S", "2011IBVS.5960....1D", "2012A&A...548A..79A", "2012IBVS.6029....1D", "2013AJ....146...38Q", "2013IBVS.6042....1D" ]	[ "10.1088/1674-4527/14/9/006", "10.48550/arXiv.1403.7954" ]	1403	1403.7954_arXiv.txt	% \label{sect:intro} The K-type short-period contact binaries, especially those with periods shorter than 0.3 days, are most probably main-sequence objects and mostly convective \citep{Bradstreet85}. They belong to W UMa type binaries which both components share common envelope that lying between the inner and outer critical Roche-lobe surfaces. W UMa type binaries have been divided into two sub-groups: A-type and W-type according to \citet{Binnendijk70}. Generally speaking, the A-type systems are more likely to have earlier spectral type than the W-type systems \citep{Rucinski74}. So it is supposed that a large amount of K-type short-period contact binaries may be W-type systems. The majority of W-type contact binaries show shallow contact characteristics (\citealt{ZhuL10}; He 2009 PhD). Therefore, they are good targets for testing the thermal relaxation oscillation theory (TRO theory; eg., \citealt{Lucy76,Flannery76,Robertson77}) V1799 Orion, or V1799 Ori (=GSC 00096-00175, $\alpha_{2000.0}$ = $04^{h}47^{m}18^{s}.19$ and $\delta_{2000.0}$ = $+06^{\circ}40'56''.1$) is a very short period (with period shorter than 0.3 days) eclipsing binary. It was first suspected to show variability by \citet{Hanley40Shap} (the cross identification is HV 10397). However, it had been neglected for a long time until ROSTE survey \citep{Akerlof00} rediscovered it to be a eclipsing binary with EW type light curves (Khruslove \citealt{Khrus1ov06}) (the cross identification is NSV 1719). The following ephemeris was reported (Khruslove \citealt{Khrus1ov06}): \begin{equation} \mathrm{Min.I (HJD)} = 2451524.829+0^{\mathrm{d}}.29031\times E. \end{equation} It was then monitored many times by several researchers such as \citet[][etc.]{Diethelm09,Diethelm10} and \citet{Nelson10}. Recently, The O-C gateway (http://var.astro.cz/ocgate/) gave out a more exact ephemeris, \begin{equation}\label{eq:ephe} \mathrm{Min.I (HJD)} = 2451524.829+0^{\mathrm{d}}.290304\times E. \end{equation} The first photometric analysis of V1799 Ori was given by \cite{Samec10} as a student/professional collaborative program. Their photometric solutions depicted V1799 Ori as a W-type, extremely shallow-contact, eclipsing binary. The degree of contact factor $3\,\%$ is rather exceptional. Their solutions uncovered two hot spots on the components, which indicates the system is quite active at present. In this paper, we present the analysis of newly obtained complete CCD light curves and compare the results with those given by Samec et al. Especially, we carry out a more exact period investigation of this system, using all the available times of light minimum.	\label{sect:discussion} The deep eclipses in both primary and secondary light minimum which denotes high inclination of the binary system help us to give reliable solutions of the system. Based on the complete BV(RI)$_{\rm c}$ light curves, the photometric solutions for V1799 Ori were carefully derived. We found that V1799 Ori is a extreme-shallow contact binary system with a mass ratio of $q = M_{2}/M_{1} = 1.335\pm0.005$ and a degree of contact about $f = 3.5\%\pm1.1\%$. It is a W-type system of which the less massive component is about 220 K hotter than the more massive one. The asymmetric light curves can be well modeled by employing a hot spot on the primary. Assuming that the primary component is a K0 type main-sequence star, its mass was roughly estimated to be $M_{2}=0.80$ M$_{\odot}$ \citep{Cox00}. Then the mass the other component was estimated to be $M_{1}=0.60$ M$_{\odot}$ by using the derived mass ratio. The photometric solution is in good agreement with that given by \cite{Samec10} except for the configuration of the star spots. This can be explained by the activity of the components. The solar-like activities on the surface of the binaries are expected to change with time, which would cause different patterns of star spots. When examining the light curves obtained by Samec et al., we found their shapes are slightly different from ours. The difference between the two maximum in their light curves are about 0.03 mag. (B,V), 0.02 mag (R), and 0.01 mag. (I), which are distinctively smaller than that in the newly obtained light curves (see Table \ref{tabmode}). This further confirms that the binary is highly active at present. Based on the analysis of the O-C diagram (Figure \ref{fig:oc}), a general trend of long-term period increase at a rate of $1.8(\pm0.6)\times10^{-8}$ days$\cdot$yr$^{-1}$ was derived. Although the period increase is insignificant (close to the error), it is no doubt a small one compared with other shallow contact binaries (see Table 7 in the paper \citealt{ZhuL10}). This is in accord with the exceptional low degree of contact, which also agrees that the period-increased system usually have a lower degree of contact \citep{ZhuL10}. The long-term period increase, together with the exceptional low degree of contact, suggests that the binary may be at a critical stage which is predicted by the TRO theory. If the period variation is caused by a conservative mass transfer, then using the well-known equation \begin{equation} \frac{\dot{P}}{P}=3\frac{\dot{M}_{2}}{M_{2}}(\frac{M_{2}}{M_{1}}-1), \end{equation} the mass transfer rate is estimated to be ${\rm d}M_{2}/{\rm d}t=1.4(\pm0.5)\times10^{-8}$ M$_{\odot}\cdot$yr$^{-1}$. However, the real mass transfer rate may be quite different from this value because of the contribution of angular momentum loss (AML) process \citep{Rahunen81,QianS13}. In addition, it is possible that the long-term increase may be only one part of a long cyclic oscillation. More observations are needed to clarify the nature of the period variation. \normalem	14	3	1403.7954	New multi-color light curves of the very short period K-type eclipsing binary V1799 Ori were obtained and analyzed with the Wilson-Devinney code. The photometric solutions reveal that the system is a W-type shallow-contact binary with a mass ratio of q = 1.335(±0.005) and a degree of contact of about f = 3.5(±1.1)%. In general, the results are in good agreement with what is reported by Samec. Dramatic manifestations of the O'Connell effect that appear in the light curves can be explained well by employing starspots on the binary surface, which confirms that the system is active at present. Several new times of light minimum were obtained. All the available times of light minimum were collected, along with the recalculated and newly obtained values. Applying a least-squares method to the constructed O — C diagram, a new ephemeris is derived for V1799 Ori. The orbital period is found to show a continuous weak increase at a rate of 1.8(±0.6) × 10<SUP>-8</SUP> d yr<SUP>-1</SUP>. The extremely shallow contact, together with the period increase, suggests that the binary may be at a critical stage predicted by thermal relaxation oscillation theory.	false	[ "thermal relaxation oscillation theory", "V1799 Ori", "New multi-color light curves", "light minimum", "K-type eclipsing binary V1799 Ori", "contact", "present", "3.5(±1.1)%", "C diagram", "Several new times", "Samec", "f", "the Wilson-Devinney code", "good agreement", "starspots", "ephemeris", "a continuous weak increase", "the light curves", "the binary surface", "a W-type shallow-contact binary" ]	6.842063	11.066401	81
485545	[ "Townsley, Leisa K.", "Broos, Patrick S.", "Garmire, Gordon P.", "Bouwman, Jeroen", "Povich, Matthew S.", "Feigelson, Eric D.", "Getman, Konstantin V.", "Kuhn, Michael A." ]	2014ApJS..213....1T	[ "The Massive Star-Forming Regions Omnibus X-Ray Catalog" ]	75	[ "Department of Astronomy & Astrophysics, 525 Davey Laboratory, Pennsylvania State University, University Park, PA 16802, USA", "Department of Astronomy & Astrophysics, 525 Davey Laboratory, Pennsylvania State University, University Park, PA 16802, USA", "Huntingdon Institute for X-ray Astronomy, LLC, 10677 Franks Road, Huntingdon, PA 16652, USA", "Max Planck Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany", "California State Polytechnic University, 3801 West Temple Ave., Pomona, CA 91768, USA", "Department of Astronomy & Astrophysics, 525 Davey Laboratory, Pennsylvania State University, University Park, PA 16802, USA", "Department of Astronomy & Astrophysics, 525 Davey Laboratory, Pennsylvania State University, University Park, PA 16802, USA", "Department of Astronomy & Astrophysics, 525 Davey Laboratory, Pennsylvania State University, University Park, PA 16802, USA" ]	[ "2014ApJ...787..107K", "2014ApJ...787..108G", "2015A&A...573A..95M", "2015A&A...574A..13E", "2015A&A...579A.131C", "2015ApJ...802...60K", "2015ApJ...811...10R", "2015ApJ...812..131K", "2015ApJ...813...42K", "2016A&A...586A.114M", "2016A&A...587A.135R", "2016A&A...594A..82R", "2016A&A...595A..27G", "2016ApJ...823...96R", "2016ApJ...825..125P", "2016ApJ...832..187B", "2016ApJ...833..193R", "2016JHEAp..11....1H", "2017A&A...598A..85S", "2017A&A...607A..86R", "2017AJ....154...87K", "2017ApJ...838...61P", "2017ApJ...850..195E", "2017ApJS..229...28G", "2017ApJS..230....3S", "2017MNRAS.468.2684P", "2017MNRAS.470.2283W", "2017arXiv170206627G", "2017arXiv171111101K", "2018A&A...615A.148D", "2018A&A...617A..14P", "2018ASPC..517..293B", "2018ASPC..517..321H", "2018ASSL..424..119F", "2018ApJ...853..171G", "2018ApJ...864..136B", "2018ApJ...865...65C", "2018ApJS..235...43T", "2018MNRAS.474.3228P", "2018MNRAS.477.5191R", "2018MNRAS.478.1218T", "2018arXiv180606293M", "2019A&A...622A..48R", "2019AJ....157...29H", "2019ApJ...878..111H", "2019ApJ...881...37P", "2019ApJS..244...28T", "2019MNRAS.482L.102R", "2019MNRAS.484.2692T", "2019MNRAS.488.1141J", "2019MNRAS.489.4278G", "2020A&A...633A.155R", "2020Ap&SS.365....6K", "2020ApJ...888..118M", "2020ApJ...897..131S", "2020ApJ...899...94R", "2020MNRAS.495.3041T", "2020PASP..132j4301S", "2021A&A...650A.156D", "2021ApJ...916...32G", "2021ApJ...921..165S", "2021PASJ...73S.129K", "2022A&A...661A..38P", "2022ApJ...933...60D", "2022ApJ...935..171B", "2022ApJ...937...46K", "2022MNRAS.510.6133B", "2022MNRAS.515.1815P", "2022MNRAS.515.4130C", "2023ApJ...958L..30R", "2023MNRAS.522.4674G", "2024A&A...682A..49G", "2024MNRAS.528..209P", "2024arXiv240316944T", "2024arXiv240509184S" ]	[ "astronomy" ]	11	[ "ISM: individual objects: NGC 6334 NGC 6357 M16 M17 W3 W4 NGC 3576 G333.6-0.2 W51A G29.96-0.02 NGC 3603 30 Doradus", "open clusters and associations: individual: Pismis 24 AH03 J1725-34.4 NGC 6611 NGC 6618 W3 Main W3(OH) IC 1795 IC 1805 OCl 352 G49.5-0.4 R136 NGC 2060", "stars: early-type", "stars: formation", "stars: individual: Pismis 24-1 Pismis 24-17 WR 93 [N78] 49 HD 168076 NGC 6611 213 CEN1a CEN1b Cl* NGC 6618 Sch 1 HD 15558 EM Car W51 IRS2E NGC 3603-A1 NGC 3603-B NGC 3603-C Cl* NGC 3603 Sher 47 WR 42e MTT 58 MTT 68 Mk34 R140a1a2", "X-rays: stars", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1980A&A....91..186C", "1984A&A...140...24L", "1989ApJS...69..831W", "1990ApJ...365..510C", "1991A&A...247..529R", "1991ApJ...370..541W", "1994A&A...282..474P", "1994ApJ...436..705G", "1994ApJS...91..713M", "1996A&A...307..775M", "1996Natur.380..687N", "1997A&A...318..931T", "1997A&A...323..121B", "1998ApJ...499L.179M", "1998ApJ...502..265H", "1999A&AS..136..313S", "1999ApJ...516..843B", "2000A&AS..143...23O", "2000ApJ...534L.139T", "2000ApJ...543..799O", "2001ApJ...558L.101R", "2002A&A...389..871D", "2002ApJ...573..191M", "2002NIMPA.486..716T", "2002NIMPA.486..751T", "2003A&A...404..223B", "2003ASPC..295..489J", "2003ApJ...590..306D", "2003ApJ...590..906T", "2003ApJ...593..874T", "2003ApJ...599.1196K", "2003PASP..115..953B", "2003SPIE.4851...28G", "2004A&A...422L..29M", "2004AJ....127.1014S", "2004AJ....127.2826B", "2004AJ....128.2374S", "2004ApJ...606..943S", "2005A&A...429..497R", "2005A&A...435...95B", "2005AJ....129..393O", "2005AJ....129.1523F", "2005ApJ...628..986G", "2005ApJS..160..319G", "2005ApJS..160..401P", "2005IAUS..227..297T", "2005MNRAS.356..801F", "2005yCat.1297....0Z", "2006A&A...456.1121D", "2006AJ....131.1163S", "2006AJ....131.2140T", "2006AJ....131.2164T", "2006ApJ...638..860E", "2006ApJ...651..237C", "2006MNRAS.367.1609B", "2006MNRAS.368.1843F", "2006SPIE.6270E..1VF", "2007ApJ...654..347L", "2007ApJ...656..227H", "2007ApJ...660..346P", "2007ApJ...660.1480M", "2007ApJ...666..321I", "2007ApJ...668..906M", "2007ApJS..168..100W", "2007ApJS..169..353B", "2007MNRAS.377..491B", "2007MNRAS.379..663M", "2007MNRAS.379.1302B", "2008A&A...490L..27A", "2008AJ....135..878M", "2008AJ....136.2083M", "2008ApJ...672.1006E", "2008ApJ...673..354F", "2008ApJ...675.1319H", "2008ApJ...686..310H", "2008MNRAS.386.1069W", "2008Sci...319..309G", "2008flhs.book.....C", "2008hsf2.book.....R", "2009AJ....138..227F", "2009ApJ...691.1109L", "2009ApJ...693..186L", "2009ApJ...693..413X", "2009ApJ...696.1278P", "2009PASP..121...76C", "2010A&A...515A..55R", "2010A&A...521A..61G", "2010A&A...523A...6D", "2010ApJ...708.1760G", "2010ApJ...714.1582B", "2010ApJ...716L..90R", "2010ApJ...719L.185J", "2010ApJ...720.1108B", "2010ApJ...725.2485K", "2010MNRAS.403.1657P", "2010MmSAI..81..171P", "2011A&A...525A...8R", "2011A&A...526A.151R", "2011A&A...530A.108E", "2011A&A...531A..51F", "2011AJ....142..158A", "2011ApJ...726...19W", "2011ApJ...731...91L", "2011ApJ...733..113R", "2011ApJ...738...34P", "2011ApJ...743...39R", "2011ApJS..194....1T", "2011ApJS..194....2B", "2011ApJS..194....4B", "2011ApJS..194....9F", "2011ApJS..194...15T", "2011ApJS..194...16T", "2011MNRAS.415.2844C", "2012A&A...538A.142R", "2012A&A...545L...2H", "2012Ap&SS.340..263C", "2012ApJ...750L..44K", "2012ApJ...753..117G", "2012ApJ...754L..37S", "2012MNRAS.427L..65R", "2012ascl.soft03001B", "2012ascl.soft12002G", "2013A&A...552A.123B", "2013A&A...554A..42R", "2013A&A...554L...2Z", "2013ApJ...764...73P", "2013ApJ...766...85R", "2013ApJS..209...26F", "2013ApJS..209...27K", "2013ApJS..209...28K", "2013ApJS..209...29K", "2013ApJS..209...30N", "2013ApJS..209...31P", "2013ApJS..209...32B", "2013MNRAS.429..398P", "2013MNRAS.431.1337R", "2013MNRAS.433..712R", "2013MNRAS.435...30W", "2013MNRAS.435.3058D", "2013MNRAS.435L..73R", "2013Natur.495...76P", "2013NewA...20...42M" ]	[ "10.1088/0067-0049/213/1/1", "10.48550/arXiv.1403.2576" ]	1403	1403.2576_arXiv.txt			14	3	1403.2576	We present the Massive Star-forming Regions (MSFRs) Omnibus X-ray Catalog (MOXC), a compendium of X-ray point sources from Chandra/ACIS observations of a selection of MSFRs across the Galaxy, plus 30 Doradus in the Large Magellanic Cloud. MOXC consists of 20,623 X-ray point sources from 12 MSFRs with distances ranging from 1.7 kpc to 50 kpc. Additionally, we show the morphology of the unresolved X-ray emission that remains after the cataloged X-ray point sources are excised from the ACIS data, in the context of Spitzer and WISE observations that trace the bubbles, ionization fronts, and photon-dominated regions that characterize MSFRs. In previous work, we have found that this unresolved X-ray emission is dominated by hot plasma from massive star wind shocks. This diffuse X-ray emission is found in every MOXC MSFR, clearly demonstrating that massive star feedback (and the several-million-degree plasmas that it generates) is an integral component of MSFR physics.	false	[ "X-ray point sources", "Omnibus X-", "MSFRs", "massive star wind shocks", "MSFR physics", "the cataloged X-ray point sources", "the unresolved X-ray emission", "this unresolved X-ray emission", "1.7 kpc", "50 kpc", "This diffuse X-ray emission", "20,623 X-ray point sources", "WISE observations", "hot plasma", "ionization fronts", "MSFR", "the Large Magellanic Cloud", "MOXC", "ACIS", "Doradus" ]	10.171709	10.815333	-1
525810	[ "Marsh, David J. E.", "Grin, Daniel", "Hlozek, Renee", "Ferreira, Pedro G." ]	2014arXiv1403.4216M	[ "Tensor Detection Severely Constrains Axion Dark Matter" ]	29	[ "-", "-", "-", "-" ]	[ "2014IJMPD..2341005P", "2014JCAP...05..046K", "2014JHEP...06..037D", "2014PhLB..734...21H", "2014PhLB..734...68K", "2014PhLB..734..183I", "2014PhLB..734..354I", "2014PhLB..735...95K", "2014PhLB..735..164C", "2014PhLB..737...30B", "2014PhLB..737..178K", "2014arXiv1403.4594V", "2014arXiv1403.5243M", "2014arXiv1407.0546R", "2015EPJP..130...69M", "2015IJMPA..3030019K", "2015arXiv150309026K", "2016arXiv160705146B", "2016arXiv161002743A", "2017arXiv170202116D", "2017arXiv170303114K", "2019BAAS...51c.567G", "2019arXiv190206567D", "2019arXiv190312643M", "2019arXiv190704473A", "2019arXiv191011591S", "2021arXiv210505771R", "2022arXiv220314915A", "2024arXiv240208716I" ]	[ "astronomy", "physics" ]	10	[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]	[ "1977PhRvD..15.2738G", "1977PhRvL..38.1440P", "1978PhRvL..40..223W", "1978PhRvL..40..279W", "1981PhRvD..23..347G", "1981RvMP...53...43G", "1982PhLB..108..389L", "1982PhRvL..48.1220A", "1983PhLB..120..127P", "1983PhLB..120..133A", "1983PhLB..120..137D", "1983PhLB..125...35T", "1983PhLB..126..178A", "1983PhRvD..28.1243T", "1984PhLB..149..351W", "1985PhRvD..32.3178S", "1986PhRvD..33..889T", "1987PhR...150....1K", "1991PhRvL..66....5T", "1992PhRvD..45.3394L", "1992SvJNP..55.1063B", "1995hep.ph....5253D", "1997PhRvL..78.1861L", "2000PhRvL..85.1158H", "2002SSRv..100..153R", "2004hep.th....9059F", "2005JCAP...01..005K", "2006JHEP...06..051S", "2006PhLB..642..192A", "2006PhRvL..97m1801B", "2008LNP...741...19S", "2008LNP...741...51R", "2008PhRvD..78h3507H", "2009ApJS..180..330K", "2009JCAP...06..022H", "2009PhRvL.103e1302P", "2010JHEP...11..105S", "2010PhRvD..81l3530A", "2010PhRvD..82j3528M", "2010PhRvD..82l3508W", "2011JCAP...07..021M", "2011PhRvD..83d4026A", "2011PhRvD..83l3526M", "2012JHEP...10..146C", "2012PhRvD..85b3534C", "2012PhRvD..85j3514M", "2012PhRvD..86b3508M", "2012PhRvD..86h3535P", "2013PhLB..727..448J", "2013PhRvD..87l1701M", "2014A&A...571A..16P", "2014A&A...571A..22P", "2014JCAP...08..016S", "2014JPhCS.485a2013R", "2014MNRAS.437.2652M", "2014PhLB..728..532F", "2014PhRvD..89j3513I", "2014PhRvL.112x1101B", "2014PhRvL.113p1801A", "2014PhRvX...4b1030B", "2014arXiv1403.4594V" ]	[ "10.48550/arXiv.1403.4216" ]	1403	1403.4216_arXiv.txt			14	3	1403.4216	The recent detection of B-modes by BICEP2 has non-trivial implications for axion dark matter implied by combining the tensor interpretation with isocurvature constraints from Planck. In this paper the measurement is taken as fact, and its implications considered, though further experimental verification is required. In the simplest inflation models $r=0.2$ implies $H_I=1.1\times 10^{14}\text{ GeV}$. If the axion decay constant $f_a<H_I/2\pi$ constraints on the dark matter (DM) abundance alone rule out the QCD axion as DM for $m_a \lesssim 52\chi^{6/7}\,\mu\text{eV}$ (where $\chi>1$ accounts for theoretical uncertainty). If $f_a>H_I/2\pi$ then vacuum fluctuations of the axion field place conflicting demands on axion DM: isocurvature constraints require a DM abundance which is too small to be reached when the back reaction of fluctuations is included. High $f_a$ QCD axions are thus ruled out. Constraints on axion-like particles, as a function of their mass and DM fraction, are also considered. For heavy axions with $m_a\gtrsim 10^{-22}\text{ eV}$ we find $\Omega_a/\Omega_d\lesssim 10^{-3}$, with stronger constraints on heavier axions. Lighter axions, however, are allowed and (inflationary) model-independent constraints from the CMB temperature power spectrum and large scale structure are stronger than those implied by tensor modes.	false	[ "axion DM", "axion dark matter", "heavier axions", "heavy axions", "isocurvature constraints", "stronger constraints", "DM", "Lighter axions", "Constraints", "further experimental verification", "non-trivial implications", "theoretical uncertainty", "tensor modes", "large scale structure", "fluctuations", "a DM abundance", "conflicting demands", "axion-like particles", "the QCD axion", "High $f_a$ QCD axions" ]	9.530189	-2.088605	65
12404171	[ "Guo, Yan-Jun", "Dai, Shi", "Li, Zhao-Sheng", "Liu, Yuan", "Tong, Hao", "Xu, Ren-Xin" ]	2015RAA....15..525G	[ "Understanding the X-ray spectrum of anomalous X-ray pulsars and soft gamma-ray repeaters" ]	4	[ "School of Physics and State Key Laboratory of Nuclear Physics and Technology, Peking University, Beijing 100871, China", "School of Physics and State Key Laboratory of Nuclear Physics and Technology, Peking University, Beijing 100871, China", "School of Physics and State Key Laboratory of Nuclear Physics and Technology, Peking University, Beijing 100871, China", "Institute of High Energy Physics, Chinese Academy of Sciences, Beijing 100049, China", "Xinjiang Astronomical Observatory, Chinese Academy of Sciences, Urumqi 830011, China", "School of Physics and State Key Laboratory of Nuclear Physics and Technology, Peking University, Beijing 100871, China; Kavli Institute for Astronomy and Astrophysics, Peking University, Beijing 100871, China" ]	[ "2015AN....336..835T", "2015MNRAS.454.3366Z", "2016RAA....16...83L", "2021JCAP...06..036F" ]	[ "astronomy" ]	9	[ "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1986ApJ...310..261A", "1992ApJ...392L...9D", "2000ApJ...534..373C", "2001ApJ...554.1245A", "2002ApJ...574..332T", "2003A&A...411L...1W", "2003A&A...411L.131U", "2003A&A...411L.141L", "2003ApJ...596L..59X", "2005MNRAS.362..777H", "2005Natur.434.1098H", "2006ApJ...645..556K", "2006MNRAS.373L..85X", "2007A&A...476..321E", "2007AdSpR..40.1453X", "2007Ap&SS.308..109B", "2007Ap&SS.308..505R", "2007ApJ...657..967B", "2007PASJ...59S..23K", "2007PASJ...59S..35T", "2008A&ARv..15..225M", "2008ApJ...680..602F", "2010A&A...518A..46T", "2010ApJ...715..665E", "2010ApJ...722L.162E", "2010ApJ...723..100S", "2010PASJ...62..475E", "2011A&A...529A..19B", "2011IJMPE..20...15T", "2013ApJ...768..144T", "2014RAA....14..673W" ]	[ "10.1088/1674-4527/15/4/006", "10.48550/arXiv.1403.6739" ]	1403	1403.6739_arXiv.txt	% Since the discovery of radio pulsars in 1967, various kinds of pulsar-like objects have been observed, with diverse manifestations. Among them, anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are peculiar kinds of sources~\citep{Mereghetti08}. Their persistent X-ray luminosities are much higher than spin down energy, while no binary signature has been observed, thus ruling out both rotation and accretion in binary system as power source. They also show recurrent bursts in the hard X-ray/soft gamma-ray band and even giant flares, which could be highly super-Eddington. Besides, AXPs/SGRs have long spin periods clustered in the range of 2 -- 12 s, and their period derivatives are large too. Conventional models for AXPs/SGRs are magnetars~\citep{Duncan92}, isolated neutron stars with extremely strong dipole and multipole magnetic fields (higher than the quantum critical magnetic field $B_{\rm QED} = m^2_ec^3/e\hbar = 4.4 \times 10^{13}$ G). The persistent emission is powered by magnetic field decay; magnetic dipole radiation would contribute to their spin-down, and the effects of twisted magnetosphere~\citep{Thompson02} and wind braking~\citep{Tong13} have also been considered. Sudden release of magnetic energy, such as magnetic reconnection, could result in bursts or giant flares. Although magnetar model could explain some of the properties of AXPs/SGRs, it is still facing some problems arising from accumulating observations: few predictions have been confirmed yet. After all, there is no direct evidence for the existence of super-strong magnetic field. Alternative models of AXPs/SGRs are not only possible but also welcome. AXPs/SGRs are suggested to be normal-field pulsar-like stars accreting from supernova fallback disks~\citep{Chatterjee00, Alpar01}. Accretion energy could power the persistent emission, and propeller effect may account for the braking mechanism, as well as the period clustering of AXPs/SGRs. However, fallback disk models could not explain the super-Eddington bursts or giant flares. The problem could be solved if the compact star is a solid quark star~\citep{Xu03}, since the self-confined surface~\citep{Alock86} could explain the super-Eddington phenomena, and energy released during star quakes~\citep{Xu06} may be alternative power source for bursts and giant flares. Therefore, AXPs/SGRs could be quark star/fallback disk systems~\citep{Xu07, TX11}. To determine whether AXPs/SGRs are magnetars or fallback disk systems is of fundamental importance. It could help us understand the observational phenomena of AXPs/SGRs, and even give hints on the nature of pulsar-like stars, which is related to the state of cold matter at supra-nuclear density and strong interaction. In this paper, we would like to study the problem from the point of X-ray spectrum of AXPs/SGRs. AXPs/SGRs have soft spectra below 10 keV that are generally fitted by a combination of a steep power-law with photon index $\Gamma \sim 2-4$ and a blackbody with temperature $kT \sim 0.5$ keV~\citep{Mereghetti08}. Non-thermal hard X-ray components above 15 keV in AXPs/SGRs were discovered in recent years, with different spectral properties from the soft X-ray band~\citep{Kuiper06}. The hard X-ray spectra are well fitted by flat power-laws with photon index $\Gamma \sim 0.5-1.5$, and the luminosity is similar to that of the soft X-ray band. Therefore, the hard X-rays provide us important information to understand the magnetic fields and surface properties, and could put strong constraints on the theoretical modeling of AXPs/SGRs. The physical mechanism of hard X-ray emission is still unknown, but some possibilities have been proposed trying to explain it. In the frame of magnetars, quantum electrodynamics model~\citep{Heyl05}, bremsstrahlung model~\citep{Beloborodov07} and resonant inverse Compton scattering model~\citep{Baring07} have been explored, predicting power-law spectra with different cutoff properties. The spectral cutoff properties of AXPs/SGRs are not well understood yet. Upper limits in MeV bands are obtained with the archival CGRO COMPTEL data of four AXPs, which indicate cutoff energy below 1 MeV~\citep{Kuiper06}. Accumulating INTEGRAL IBIS data over 9 years, the time-averaged spectrum of 4U 0142+61, the brightest AXP, can be fitted with a power-law with an exponential cutoff at $\sim 130$ keV, which might rule out models involving ultra-relativistic electrons~\citep{Wang13}. In the fallback disk frame, ~\citet{Trumper10} considered producing the hard X-ray emission by bulk-motion Comptonization (BMC) process of surface photons in the accretion flow. Their work shows that for 4U 0142+61, BMC model could reproduce both the soft and hard X-ray spectra. BMC model is successful in explaining the spectra of 4U 0142+61, but its applicability to other sources remains in doubt. In this work, we try to apply BMC model to other AXPs/SGRs, and make simulation to discuss how to distinguish different models by future observations. Using data from {\it Suzaku} and INTEGRAL, we derive the soft and hard X-ray spectra of four sources, namely AXP 1RXS J170849-400910, AXP 1E 1547.0-5408, SGR 1806-20 and SGR 0501+4516 (hereafter abbreviated as 1RXS J1708-40, 1E 1547-54, SGR 1806-20 and SGR 0501+45), to put further constraints on BMC model. We find that the spectra of all the chosen sources can be well fitted with XSPEC model compTB, showing that accretion scenario could be compatible with X-ray emission of AXPs/SGRs. To investigate the feasibility of discriminating various models of hard X-ray emission by future observations, we also make simulation for the hard X-ray modulation telescope (HXMT) \footnote{http://heat.tsinghua.edu.cn/hxmtsci/hxmt.html}. Simulated spectra of BMC model exhibit cutoff around 200 keV, which could distinguish BMC from other cases in the magnetar model. In \S~2 we will introduce the utilized {\it Suzaku} and INTEGRAL observations, along with data analysis including spectral properties and time variabilities. Then we present the averaged spectra and the fitting of compTB model in \S~3. HXMT simulations are shown in \S~4, and possible discrepancies between various models of the hard tail are also discussed. Finally we make conclusions in \S~5.	Whether AXPs/SGRs are magnetars or quark star/fallback disk systems remains a problem to be settled. We study the soft and hard X-ray spectra of four AXPs/SGRs with {\it Suzaku} and INTEGRAL observations.% The broad-band {\it Suzaku} spectra could be well reproduced by BMC process, and BMC model could also fit the combined {\it Suzaku} and INTEGRAL spectra, with parameters better constrained. Thus fallback disk system could be compatible with X-ray emission of AXPs/SGRs, implying that the existence of accretion flow is possible. In addition, HXMT simulated spectra of BMC model exhibit cutoff around 200 keV, showing significant discrepancy from power-law spectra. We can expect to distinguish BMC model from other hard X-ray models in the magnetar frame by future Chinese HXMT observations, and further understand the nature of AXPs/SGRs.	14	3	1403.6739	Hard X-rays above 10 keV are detected from several anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs), and different models have been proposed to explain the physical origin within the frame of either a magnetar model or a fallback disk system. Using data from Suzaku and INTEGRAL, we study the soft and hard X-ray spectra of four AXPs/SGRs: 1RXS J170849-400910, 1E 1547.0-5408, SGR 1806-20 and SGR 0501+4516. It is found that the spectra could be well reproduced by the bulk-motion Comptonization (BMC) process as was first suggested by Trümper et al., showing that the accretion scenario could be compatible with X-ray emission from AXPs/SGRs. Simulated results from the Hard X-ray Modulation Telescope using the BMC model show that the spectra would have discrepancies from the power-law, especially the cutoff at ∼200 keV. Thus future observations will allow researchers to distinguish different models of the hard X-ray emission and will help us understand the nature of AXPs/SGRs. <P />Supported by the National Natural Science Foundation of China.	false	[ "SGR", "X-ray emission", "several anomalous X-ray pulsars", "different models", "SGRs", "rays", "-", "the hard X-ray emission", "1RXS J170849", "the Hard X-ray Modulation Telescope", "Trümper et al", "soft gamma-ray repeaters", "the soft and hard X-ray spectra", "Hard X", "1E", "AXPs/SGRs", "AXPs", "BMC", "Modulation Telescope", "China" ]	5.298963	4.164489	69
436902	[ "Loukitcheva, M.", "Solanki, S. K.", "White, S. M." ]	2014A&A...561A.133L	[ "The chromosphere above sunspots at millimeter wavelengths" ]	26	[ "Max-Planck-Institut for Sonnensystemforschung, 37191, Katlenburg-Lindau, Germany ; Astronomical Institute, St. Petersburg University, Universitetskii pr. 28, 198504, St. Petersburg, Russia", "Max-Planck-Institut for Sonnensystemforschung, 37191, Katlenburg-Lindau, Germany; School of Space Research, Kyung Hee University, Yongin, Gyeonggi, 446-701, Korea", "Space Vehicles Directorate, Air Force Research Laboratory, Kirtland AFB, NM, USA" ]	[ "2015ASPC..499..351F", "2015ApJ...804...48I", "2015CosRe..53...10B", "2016ApJ...816...91I", "2016BaltA..25..225R", "2016SSRv..200....1W", "2017A&A...601A..43L", "2017ARA&A..55..159L", "2017ApJ...841L..20I", "2017ApJ...845..102H", "2017ApJ...850...35L", "2017SoPh..292...22I", "2017SoPh..292...88W", "2018A&A...613A..17B", "2018A&A...620A.124D", "2018SoPh..293...13S", "2019AdSpR..63.1396L", "2019SoPh..294..163S", "2020Ge&Ae..59..783B", "2020SoPh..295....4R", "2021ApJ...907...16B", "2022FrASS...9.1118D", "2022FrASS...925368L", "2023arXiv231212210F", "2024arXiv240300920W", "2024arXiv240504871P" ]	[ "astronomy" ]	3	[ "Sun: chromosphere", "Sun: radio radiation", "sunspots", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1968SvA....12..245Z", "1971SoPh...18..347G", "1981ApJS...45..635V", "1981phss.conf..235A", "1986ApJ...306..284M", "1988A&A...189..232O", "1993A&A...276..219C", "1993ApJ...406..319F", "1993ApJ...415..364B", "1994ASIC..433..169S", "1994SoPh..152..175N", "1995ApJ...453..517L", "1995SoPh..162..129S", "1996ASSL..204.....Z", "1996PASP..108...93W", "1997SoPh..174...31W", "2001A&A...375..591C", "2003A&ARv..11..153S", "2004A&A...419..747L", "2006A&A...456..697W", "2006SPIE.6267E..13B", "2006SoPh..239...41L", "2007A&A...465..291M", "2007ApJ...667.1243F", "2007ApJS..169..439S", "2008ApJS..175..229A", "2009A&A...505..307T", "2009ApJ...707..482F", "2009SoPh..255...53U", "2009SoPh..260..233W", "2013A&A...557A..24V" ]	[ "10.1051/0004-6361/201321321", "10.48550/arXiv.1403.3436" ]	1403	1403.3436_arXiv.txt	There have been numerous attempts to build a comprehensive semiempirical model of the atmosphere above a sunspot umbra. Deriving such models is a complicated process that tries to balance observations of a range of optical and UV lines \citep[mostly formed in non-LTE conditions, see, e.g.,][for a review see Solanki 2003]{Avrett1981, Maltby1986, Obridko} and, when available, radio measurements of brightness spectra, together with ionization equilibrium and radiative transfer calculations that include heat transfer down from the corona as well as other factors \citep[e.g.,][]{Fontenla1993}. The radio data are particularly valuable for solar diagnostics because the measurements are in the Rayleigh--Jeans limit, meaning that measured brightness temperatures actually represent thermal electron temperatures in the optically thick atmosphere. The temperature, at which a given frequency is optically thick, is sensitive to density and temperature, and hence the radio data provide important constraints for modeling \citep[e.g.,][]{loukitcheva}. However, while there are numerous radio temperature measurements of the quiet--Sun atmosphere, there are very few for sunspots, particularly at millimeter wavelengths, since good spatial resolution is required to isolate the brightness temperature of the sunspot from its surroundings, and single--dish measurements generally do not have enough resolution. The highest--resolution single--dish data, at $10-20$\arcsec, are from JCMT \citep{lindsey} and CSO \citep{bastian}, but at sub-mm wavelengths, which sample deeper in the atmosphere, closer to photospheric heights. The sunspot umbra appears darker than the quiet Sun at these wavelengths and umbral contrast decreases towards longer wavelengths. It is more interesting, however, to use millimeter wavelength radio data, which can constrain the temperature in the chromosphere and address the controversy regarding the umbral temperature in these layers \citep[see, e.g.,][and Sect.~\ref{section3}]{Maltby1986, Severino, Fontenla2009}. At millimeter wavelengths the instrument best suited for high-resolution observations of sunspots until recently was the 10-element Berkeley-Illinois-Maryland Array (BIMA) operating at $3.5$~mm \citep{welch}. The BIMA antennas are now part of the Combined Array for Millimeter Astronomy \citep[CARMA;][]{Bock2006}. \begin{figure}[!htb] \centering \includegraphics[width=0.32\textwidth, angle=90]{fig_maps_paper_upd3.eps} \caption{Active region NOAA 10448 observed on August 31, 2003. \textit{(a)} BIMA image at 3.5 mm, \textit{(b)} MDI photospheric magnetogram and \textit{(c)} BBSO Ca~{\sc ii}~K image. White contours mark sunspot umbrae and penumbrae. The box contains the region shown in Fig.~\ref{fig3} as a blowup. The zero brightness in the BIMA image corresponds to the weakest flux.} \label{fig1} \end{figure} \begin{figure}[!htb] \centering \includegraphics[width=0.25\textwidth, angle=90]{fig_maps_paper_over_upd_new.eps} \caption{BBSO Ca~{\sc ii}~K image and MDI photospheric magnetogram with the overlaid contours of 3.5 mm emission plotted as iso-intensity lines corresponding to 10, 20, 30, 40, 50, 60, 70 \% of maximum brightness. } \label{fig22} \end{figure} \begin{figure}[!htb] \centering \includegraphics[width=0.25\textwidth, angle=90]{fig_qs_40G_paper_upd.eps} \caption{\textit{(a)} MDI magnetogram in the range (-200,200)~Gauss. The dotted and solid lines mark the 40~G and 26~G contours, corresponding to the quiet-Sun level at the original resolution of the magnetogram and at the resolution of the mm data, respectively. The white circle corresponds to the BIMA primary beam of 3\arcmin\ by 3\arcmin\ containing flux.\textit{(b)} The resulting quiet-Sun mask restricted to BIMA's primary beam. Quiet-Sun areas are white. } \label{fig2} \end{figure} In this work we use BIMA observations of the temperature contrast (relative to the quiet Sun) above a sunspot umbra at $3.5$~mm. At this wavelength models predict rather different umbral brightnesses, so that such measurements can be used to distinguish between the various atmospheric models of sunspots. In Section~\ref{section2} we report the observations of the active region at $3.5$~mm, evaluate umbral brightness at this wavelength relative to the quiet Sun, and investigate the influence of spatial resolution on the appearance of the sunspot at mm wavelengths. In Section~\ref{section3} we discuss the differences between the temperature and density stratifications in the existing sunspot models as well as between mm brightness spectra calculated from these models. We demonstrate that observations of sunspots at mm wavelengths impose strong constraints on temperature and density stratifications of the sunspot atmosphere, particularly on the location and depth of the temperature minimum and the location of the transition region in the umbral models. In Section~\ref{section4} we complement the investigation of sub-mm and mm umbral brightnesses with two examples of the observations of sunspot umbra at short cm wavelengths.	Millimeter brightness observations impose strong constraints on temperature and density stratifications of the sunspot atmosphere, in particular on the location and depth of the temperature minimum and the location of the transition region. Current submm/mm observational data suggest that existing sunspot models, including the models of \citet{Avrett1981}, \citet{Maltby1986}, \citet{Obridko}, \citet{socas} and \citet{Fontenla2009}, fail to reproduce the millimeter observations and therefore are incomplete descriptions of the umbral atmosphere from the temperature minimum through the chromosphere up to the transition region. Only the model of \citet{Severino} lies within the error bars at all the wavelengths considered here. From the analysis of the existing models we can conclude that a successful model that is in agreement with submillimeter and millimeter umbral brightness should have an extended or/and deep temperature minimum (3000~K or below) as, for instance, in the models of \citet{Severino} and \citet{Fontenla2009}. % Most atmospheric models are based on one-dimensional static atmospheres that try to reproduce observed umbral spectra. However, according to J.~Fontenla (priv. comm.) there is an obvious reason why such 1D models will have difficulty. The atmosphere of a sunspot umbra is strongly irradiated from the sides by the penumbra in the walls produced by the Wilson depression, and this will affect the observed spectra. An example is the fact that Lyman-$\alpha$ has centrally--peaked profiles in the umbra but centrally--reversed profiles elsewhere \citep[e.g.,][]{Tian09}: centrally--peaked profiles are expected if the umbral Ly$\alpha$ emission is dominated by scattered and redistributed emission coming in from the penumbral walls, rather than emission intrinsic to the umbral atmosphere. One-dimensional models cannot handle the complication introduced by scattering from a source with a different atmospheric structure. This implies that umbra is probably cooler and less dense than the current 1D models indicate, which is consistent with the fact that the radio data show the depression continuing to longer wavelengths than the 1D models suggest. Due to the spatial resolution limit of 12\arcsec\ in the BIMA observations used here we are not able to fully resolve the umbra and cleanly separate it from penumbra. Therefore the results obtained in this work are preliminary. A detailed study of the appearance of sunspot umbrae at mm waves requires significantly higher spatial resolution. Furthermore, good wavelength coverage is needed for accurate diagnostics of the turnover wavelength, which is, in turn, required for the successful modeling of the sunspot atmospheric temperature structure based on mm-wavelength data. We place high expectations on the Atacama Large Millimeter/Submillimeter Array (ALMA), which has commenced Early Science observations. The current instrument wavelength range covers 0.4 to 3.6~mm and the field of view ranges from 8.5\arcsec\ at the shortest wavelengths to 72\arcsec\ at the longest wavelength. At the shortest wavelengths the single pointing field of view will be too small for umbral observations and the mosaicing observing mode will be essential to cover a sunspot together with surrounding quiet-Sun areas. With sub-arcsecond resolution and up to 66 antennas, ALMA will be an extraordinarily powerful instrument for studying the three-dimensional thermal structure of sunspots at chromospheric heights.	14	3	1403.3436	<BR /> Aims: The aim of this paper is to demonstrate that millimeter wave data can be used to distinguish between various atmospheric models of sunspots, whose temperature structure in the upper photosphere and chromosphere has been the source of some controversy. <BR /> Methods: We use observations of the temperature contrast (relative to the quiet Sun) above a sunspot umbra at 3.5 mm obtained with the Berkeley-Illinois-Maryland Array (BIMA), complemented by submm observations from Lindsey & Kopp (1995) and 2 cm observations with the Very Large Array. These are compared with the umbral contrast calculated from various atmospheric models of sunspots. <BR /> Results: Current mm and submm observational data suggest that the brightness observed at these wavelengths is low compared to the most widely used sunspot models. These data impose strong constraints on the temperature and density stratifications of the sunspot umbral atmosphere, in particular on the location and depth of the temperature minimum and the location of the transition region. <BR /> Conclusions: A successful model that is in agreement with millimeter umbral brightness should have an extended and deep temperature minimum (below 3000 K). Better spatial resolution as well as better wavelength coverage are needed for a more complete determination of the chromospheric temperature stratification above sunspot umbrae.	false	[ "submm observations", "sunspot umbrae", "sunspots", "various atmospheric models", "observations", "millimeter umbral brightness", "the sunspot umbral atmosphere", "millimeter wave data", "BIMA", "the chromospheric temperature stratification", "better wavelength coverage", "the temperature minimum", "Lindsey", "Current mm and submm observational data", "the temperature contrast", "Kopp", "amp", "whose temperature structure", "a sunspot umbra", "Berkeley" ]	12.068113	15.191821	2
2775166	[ "Bertin, E.", "Pillay, R.", "Marmo, C." ]	2015A&C....10...43B	[ "Web-based visualization of very large scientific astronomy imagery" ]	14	[ "Univ Pierre et Marie Curie, Institut d'Astrophysique de Paris, UMR7095, Paris, F-75014, France; CNRS, France", "C2RMF, Palais du Louvre - Porte des Lions, Paris 75001, France", "Univ Paris-Sud, Laboratoire GEOPS, UMR8148, Orsay, F-91405, France; CNRS, France" ]	[ "2015A&A...577A.148B", "2015A&C....13...12N", "2015A&C....13...50M", "2017A&C....20..128Y", "2017PASA...34...23M", "2017PASP..129e8006B", "2018ApJS..239...18A", "2019ASPC..521..651B", "2019JPhCS1167a2075N", "2020MNRAS.498.2138B", "2020NCimR..43..515S", "2021A&A...647A.100S", "2021ApJS..255...20A", "2022A&C....3900586H" ]	[ "astronomy", "general" ]	17	[ "Visualization", "Scientific data", "Web application", "High resolution", "HTML5", "Astrophysics - Instrumentation and Methods for Astrophysics", "Computer Science - Computational Engineering", "Finance", "and Science", "Computer Science - Multimedia", "85-08", "J.2", "I.4.9" ]	[ "1981A&AS...44..363W", "1981IEEEP..69...14H", "1996A&AS..117..393B", "2000A&AS..143...23O", "2002A&A...395.1077C", "2009PASP..121..414P", "2012ApJS..203...21A" ]	[ "10.1016/j.ascom.2014.12.006", "10.48550/arXiv.1403.6025" ]	1403	1403.6025_arXiv.txt	\label{chap:intro} Although much of the extraction of information from astronomy science images is now performed ``blindly'' using computer programs, astronomers still rely on visual examination for a number of tasks. Such tasks include image quality control, assessment of morphological features, and debugging of measurement algorithms. The generalization of standardized file formats in the astronomy community, such as FITS \citep{wells_fits_1981}, has facilitated the development of universal visualization tools. In particular, {\sc SAOimage} \citep{1990BAAS...22..935V}, {\sc Aladin} \citep{1994ASPC...61..215B}, {\sc SkyCat} \citep{1997ASPC..125..333A}, {\sc Gaia} \citep{2000ASPC..216..615D} and {\sc ds9} \citep{2003ASPC..295..489J}. These packages are designed to operate on locally stored data and provide efficient access to remote image databases by downloading sections of FITS data which are subsequently read and processed locally for display; all the workload, including image scaling, dynamic range compression, color compositing and gamma correction, is carried out client-side. However, the increasing gap between storage capacities and data access bandwidth \citep{budman} makes it increasingly efficient to offload part of the image processing and data manipulations to the server, and to transmit some form of pre-processed data to clients over the network. Thanks to the development of wireless networks and light mobile computing (tablet computers, smartphones), more and more scientific activities are now being carried out on-the-go outside an office environment. These possibilities are exploited by an increasing number of scientists, especially experimentalists involved in large international collaborations and who must interact remotely, often in real-time, with colleagues and data located in different parts of the world and in different time zones. Mobile devices have increasingly improved display and interfacing capabilities, however, they offer limited I/O performance and storage capacity, as well as poor battery life when under load. Web-based clients, or simply {\em Web Apps}, are the applications of choice for these devices, and their popularity has exploded over the past few years. Thanks to the ubiquity of web browsers on both desktop and mobile platforms, {\it Web Apps} have become an attractive solution for implementing visual interfaces. Modern web browsers feature ever faster and more efficient JavaScript engines, support for advanced standards such as HTML5 \citep{html5} and CSS3 \citep{css}, not to mention interactive 3D-graphics with the recent WebGL API \citep{webgl}. As far as data visualization is concerned, web applications can now be made sufficiently feature-rich so as to be able to match many of the functions of standalone desktop applications, with the additional benefit of having instant access to the latest data and being embeddable within web sites or data portals. One of the difficulties in having the browser deal with science data is that browser engines are designed to display gamma-encoded images in the GIF, JPEG or PNG format, with 8-bits per Red/Green/Blue component, whereas scientific images typically require linearly quantized 16-bit or floating point values. One possibility is to convert the original science data within the browser using JavaScript, either directly from FITS \citep{jsfits, astrojs}, or from a more ``browser-friendly'' format, such as e.g., a special PNG ``representation file'' \citep{js9}, or compressed JSON \citep{2011ASPC..442..467F}. In practice this is currently limited to small rasters, as managing millions of such pixels in JavaScript is still too burdensome for less powerful devices. Moreover, lossless compression of scientific images is generally not very efficient, especially for noisy floating-point data \citep[e.g.][]{2009PASP..121..414P}. Hence, currently, server-side compression and encoding of the original data to a browser-friendly format remains necessary in order to achieve a satisfactory user experience on the web client, especially with high resolution screens. Displaying images larger than a few megapixels on monitors or device screens requires panning and/or pixel rebinning, such as in ``slippy map'' implementations (Google Maps\texttrademark, OpenStreetMap\footnote{\url{http://www.openstreetmap.org}} etc.). On the server, the images are first decomposed into many small tiles (typically $256\times 256$ pixels) and saved as PNG or JPEG files at various levels of rebinning, to form a ``tiled pyramid''. Each of these small files corresponds to a URL and can be loaded on demand by the web client. Notable examples of professional astronomy web apps based on this concept include the Aladin Lite API \citep{2012ASPC..461..443S}, and the Mizar plugin\footnote{\url{https://github.com/TPZF/RTWeb3D}} in SITools2 \citep{2012ASPC..461..821M}. However, having the data stored as static 8-bit compressed images means that interaction with the pixels is essentially limited to passive visualization, with little latitude for image adjustment or interactive analysis. Server-side dynamic processing/conversion of science-data on the server and streaming of the processed data to the web client are necessary to alleviate these limitations. Visualization projects featuring dynamic image conversion/streaming in Astronomy or Planetary Science have mostly relied on browser plugins implementing proprietary technologies \citep{2012ASPC..461...95F} or Java clients/applets \citep{10.1109/MCSE.2009.142, kitaeff_2012}. Notable exceptions include Helioviewer \citep{hughitt_helioviewer:_2008}, which queries compressed PNG tiles directly from the browser with the tiles generated on-the-fly server-side from JPEG2000 encoded data. In this paper we describe an open source and multi-platform high performance client-server system for the processing, streaming and visualization of full bit depth scientific imagery at the terabyte scale. The system consists of a light-weight C++ server and W3C standards-based JavaScript clients capable of running on stock browsers. In section \ref{chap:method}, we present our approach, the protocols and the implementation of both the server and the client. Sections \ref{chap:astrapp} and \ref{chap:planetapp} showcase several applications in Astronomy and Planetary Science. In Section \ref{chap:perf}, we assess the performance of the system with various configurations and load patterns. Finally in Section \ref{chap:conclu}, we discuss future directions in the light of current technological trends.	\label{chap:conclu} A high performance web-based system for remote visualization of full resolution scientific grade astronomy images and data has been developed. The system is entirely open-source and capable of efficiently handling full resolution 32 bit floating point image and elevation map data. We have studied the performance and scalability of the system and have shown that it is capable of handling terabyte-size scientific-grade images that can be browsed comfortably by at least a hundred simultaneous users, on a single server. By using and extending an existing open source project, a system for astronomy has been put together that is fully mature, that will benefit from the synergies of the wider scientific imaging community and that is ready for use in a busy production environment. In addition the {\sc IIPImage} server, is distributed as part of the default Debian, Ubuntu and Fedora package repositories, making installation and configuration of the system very straightforward. All the code developed within this project for {\tt iipsrv} has been integrated into the main code base and will form an integral part of the 1.0 release. However, there are still many potential directions for improvements, both server-side and client-side. Most importantly: \begin{itemize} \item {The TIFF storage format used on the server currently restricts pixel bit depth, the number of image channels, and I/O performance (through {\tt libtiff}). A valid alternative to TIFF could be the Hierarchical Data Format Version 5 (HDF5) \citep{the_hdf_group_hierarchical_2000}, which provides a generic, abstract data model that enables POSIX-style access to data objects organized hierarchically within a single file; some radio-astronomers have been trying to promote the use of HDF5 for storing massive and complex astronomy datasets \citep{2012ASPC..461..871M}. A more radical approach would be to adopt JPEG2000 as the archival storage format for astronomy imaging archives \citep{2014arXiv1403.2801K}, which could also remove the need for transcoding images for visualization purposes.} \item {Additional image operations could be implemented within {\tt iipsrv}, including real-time hyperspectral image processing and compositing.} \item {Although the {\sc IIPImage} image tile server already supports simple standard tile query protocols and interfaces easily with most image panning clients, a welcome addition would be to offer support for the more GIS-oriented WTMS (Web Map Tile Service) protocol \citep{wmts}}. \item {The International Virtual Observatory Alliance (IVOA) has agreed on a standard set of specifications for discovering and accessing remote astronomical image datasets: the Simple Image Access Protocol (SIAP) \citep{2011arXiv1110.0499T}. The response to an SIAP query consists of metadata and download URLs for matching image products. Current SIAP specifications\footnote{http://www.ivoa.net/documents/SIA/} do not provide specific ways to access pyramids of tiled images. Still, support for SIAP could be implemented within or outside of {\tt iipsrv} for generating, for example, JPEG cutouts or lists of tiles that match a given set of coordinates/field of view/pixel scale.} \item {Both {\sc IIPMooViewer} and {\sc Leaflet} clients require all layers displayed on a map at the same moment to share the same ``native'' pixel grid (projection). Although this limitation does not prevent ``blinking'' images with different pixel grids, it precludes overlapping different observations/exposures on screen. For instance it makes it impossible to display accurately the entire focal plane of a mosaic camera on a common viewport, without prior resampling. Having different images with different native pixel grids sharing the same map would require the web-client to perform real-time reprojection. Client-side reprojection should be possible e.g., with version 3 of the {\sc OpenLayers} library\footnote{\url{http://ol3js.org/}}}. \end{itemize}	14	3	1403.6025	Visualizing and navigating through large astronomy images from a remote location with current astronomy display tools can be a frustrating experience in terms of speed and ergonomics, especially on mobile devices. In this paper, we present a high performance, versatile and robust client-server system for remote visualization and analysis of extremely large scientific images. Applications of this work include survey image quality control, interactive data query and exploration, citizen science, as well as public outreach. The proposed software is entirely open source and is designed to be generic and applicable to a variety of datasets. It provides access to floating point data at terabyte scales, with the ability to precisely adjust image settings in real-time. The proposed clients are light-weight, platform-independent web applications built on standard HTML5 web technologies and compatible with both touch and mouse-based devices. We put the system to the test and assess the performance of the system and show that a single server can comfortably handle more than a hundred simultaneous users accessing full precision 32 bit astronomy data.	false	[ "large astronomy images", "mobile devices", "image settings", "current astronomy display tools", "survey image quality control", "public outreach", "interactive data query", "standard HTML5 web technologies", "citizen science", "floating point data", "remote visualization", "terabyte scales", "full precision", "datasets", "32 bit astronomy data", "extremely large scientific images", "ergonomics", "speed", "terms", "exploration" ]	10.403322	4.52518	170
483849	[ "Takami, Hajime", "Nozawa, Takaya", "Ioka, Kunihito" ]	2014ApJ...789L...6T	[ "Dust Formation in Macronovae" ]	26	[ "Institute of Particle and Nuclear Studies, KEK, 1-1, Oho, Tsukuba 305-0801, Japan ; JSPS Research Fellow.;", "National Astronomical Observatory of Japan, 2-21-1, Osawa, Mitaka, Tokyo 181-8588, Japan", "Institute of Particle and Nuclear Studies, KEK, 1-1, Oho, Tsukuba 305-0801, Japan ; Department of Particle and Nuclear Physics, the Graduate University for Advanced Studies, Tsukuba 305-0801, Japan;" ]	[ "2015ApJ...800...79W", "2015ApJ...802..119K", "2015ApJ...804L..16K", "2015ApJ...809L...8K", "2015ApJ...811....4K", "2015IJMPD..2430012R", "2015PhRvD..92d4028K", "2015PhRvD..92f4034K", "2016AdAst2016E...8T", "2016ApJ...818..104K", "2016ApJ...823...15W", "2016ApJ...827...83K", "2017ApJ...849L..19G", "2017LRR....20....3M", "2017RPPh...80h4901K", "2018ApJ...861...55M", "2018ApJ...862L..11V", "2018PTEP.2018d3E02I", "2019LRR....23....1M", "2020MNRAS.496.1891B", "2021ApJ...923..219W", "2022A&A...666A..67P", "2022ApJ...939...59A", "2022ApJ...941..100S", "2022MNRAS.515L..89H", "2023ARNPS..73..365F" ]	[ "astronomy" ]	3	[ "binaries: general", "dust", "extinction", "gamma-ray burst: individual: GRB 130603B", "infrared: stars", "methods: numerical", "stars: neutron", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1985JPCRD..14....1W", "1996MNRAS.282.1321Z", "1997A&A...319..122R", "1998ApJ...507L..59L", "1999A&A...341..499R", "2000ApJ...530..757P", "2001A&A...380..544R", "2003ApJ...598..785N", "2003MNRAS.345.1077R", "2005ApJ...630L.117H", "2005astro.ph.10256K", "2009ApJ...690.1681D", "2009ApJ...704..306K", "2010CQGra..27h4004K", "2010MNRAS.406.2650M", "2010NIMPA.624..223A", "2011ApJ...734L..36S", "2011ApJ...738L..32G", "2011CQGra..28k4002A", "2011Natur.478...82N", "2012MNRAS.426.1940K", "2013ApJ...773...78B", "2013ApJ...774...25K", "2013ApJ...774L..23B", "2013ApJ...775...18B", "2013ApJ...775..113T", "2013ApJ...775L..19J", "2013ApJ...776...24N", "2013ApJ...778L..16H", "2013MNRAS.430.2121P", "2013MNRAS.435..502F", "2013Natur.500..547T", "2013PhRvD..87b4001H", "2013PhRvD..88d1503K", "2014ApJ...780...31T", "2014ApJ...782L...2I", "2014ApJ...789L..39W", "2014JPhG...41d4006S", "2014MNRAS.437L...6K", "2014MNRAS.439..744R", "2014MNRAS.441.3444M", "2014PhRvD..89f3006T", "2018LRR....21....3A" ]	[ "10.1088/2041-8205/789/1/L6", "10.48550/arXiv.1403.5872" ]	1403	1403.5872_arXiv.txt	\label{sec:intro} Macronovae are brightening phenomena associated with the ejecta from the mergers of neutron star binaries (NSBs), i.e., neutron star (NS)-NS binaries and black hole (BH)-NS binaries. In the original macronova model, the luminosity peaks at $\sim 1$ day after the mergers, with the opacity coefficient of ejecta $\kappa \sim 0.1$ cm$^2$ g$^{-1}$ \citep{Li1998ApJ507L59,Kulkarni2005astro-ph0510256,Metzger2010MNRAS406p2650}. Recent studies have shown that r-process nucleosynthesis occurs efficiently in neutron-rich ejecta \citep[low electron fraction $Y_e$, e.g.,][]{Goriely2011ApJ738L32,Korobkin2012MNRAS426p1940,Bauswein2013ApJ773p78}. The r-process elements, especially lanthanoids, provide large opacity for ejecta \citep[$\kappa \sim 10$ cm$^2$ g$^{-1}$;][]{Barnes2013ApJ775p18,Tanaka2013ApJ775p113,Tanaka2014ApJ780p31}, so that their luminosity peaks around $10$ days (called the r-process model). In both cases, radioactive decay heats the ejecta and powers emission. The r-process model successfully reproduces the near-infrared (NIR) macronova in the afterglow of short gamma-ray burst (GRB) 130603B \citep[$z = 0.356$;][]{Berger2013ApJ774L23,Tanvir2013Nature500p547,Hotokezaka2013ApJ778L16}. NSB mergers are the most promising sources of gravitational waves which are expected to be directly detected by the next generation interferometers, such as advanced LIGO \citep{Abadie2010NIMPhysResSectA624p223}, advanced VIRGO \citep{Accadia2011CQGra28p114002}, and KAGRA \citep{Kuroda2010CQGra27p084004}. The electromagnetic detection of macronovae improves the localization of gravitational wave sources; the localization accuracy by photons is much better than that by the interferometers, $\sim 10$ - $100$ deg$^2$ \citep[e.g.,][]{Aasi2013arXiv1304.0670}. Recent numerical simulations have revealed that NSB mergers eject significant masses with $\sim 10^{-4} M_{\odot}$ -- $10^{-2} M_{\odot}$ dynamically \citep[e.g.,][]{Rosswog1999A&A341p499,Ruffert2001A&A380p544,Hotokezaka2013PRD87p024001,Bauswein2013ApJ773p78,Kyutoku2013PRD88p041503} and/or by neutrino-driven \citep[e.g.,][]{Ruffert1997A&A319p122,Rosswog2003MNRAS345p1077,Dessart2009ApJ690p1681,Metzger2014arXiv1402.4803} or magnetically driven winds \citep[e.g.,][]{Shibata2011ApJ734L36}, which can explain GRB 130603B in the r-process model. The ejecta also interact with circumstellar matter and radiate like supernova remnants at a later phase \citep{Nakar2011Nature478p82,Piran2013MNRAS430p2121,Takami2013arXiv1307.6805,Kyutoku2014MNRAS437L6}. Different types of nucleosynthesis may take place in the ejecta, depending on $Y_e$. While r-process nucleosynthesis occurs in low $Y_e$ ejecta, relatively high $Y_e$ ($\sim 0.2$ -- $0.5$) can be also realized, which has been exclusively discussed for neutrino-driven winds \citep[e.g.,][see also \citet{Wanajo2014arXiv1402.7317} for locally low $Y_e$ dynamical ejecta]{Fernandez2013MNRAS435p502,Rosswog2013arXiv1307.2939,Surman2013arXiv1312.1199,Metzger2014arXiv1402.4803}. In such environments, r-process nucleosynthesis is inefficient, but heavy elements (up to $^{56}$Ni) may be synthesized from the constituent nucleons of NS matter, e.g., through a series of captures of $\alpha$ particles by $^{12}$C produced by the triple-$\alpha$ process \citep[e.g.,][]{Surman2013arXiv1312.1199}. In this Letter, we investigate dust formation in the ejecta of NSB mergers for the first time. The formation of dust in macronovae is expected as in supernovae \citep[e.g.,][]{Nozawa2003ApJ598p785} because heavy elements may be synthesized and the ejecta temperature may be low enough for dust formation. We demonstrate that the newly formed dust can be responsible for the opacity of ejecta and its emission can potentially reproduce the NIR excess of GRB 130603B. Although the r-process model can explain this macronova, it is based on the limited observational data. Thus, it is worth considering the dust scenario \citep[e.g., see ][for another possibility]{Jin2013ApJ775L19}.	\label{sec:dissum} We have investigated dust formation in macronovae based on the temperature and density estimated from GRB 130603B. We have shown that dust of r-process elements hardly form even if they are abundantly produced. On the other hand, dust of light elements such as carbon can be formed. We have also suggested that the NIR macronova of GRB 130603B can be explained by the emission of light-element dust such as carbon grains, as an alternative to the r-process model. We inferred the temperature of ejecta from the observational result as $T_0 \sim 2000$ K. A heating source may be the radioactive decay products of r-process elements in the r-process model. In r-process nucleosynthesis inefficient ejecta one possibility is the radioactive decay of heavy, but not r-process, elements \citep{Barnes2013ApJ775p18}. Radioactive nuclei with the lifetime less than $\sim 10$ days release a significant fraction of radioactive energies and achieve $T \sim 2000$ K at $\sim 7$ days under a reasonable choice of ejecta mass. Shock heating may be also possible. We should keep in mind that the discussions in this Letter are based on the observational result of GRB 130603B because this is the only existing sample. For instance, if ejecta temperature is lower and density is much higher in another macronova, dust grains of r-process elements could be formed. Our results have shown that newly formed grains are relatively small, consisting of $\sim 100$ up to $\sim 10^5$ atoms. We adopted the theory of Mie scattering in calculating the absorption coefficients of dust. However, they might deviate from the prediction of the theory for the dust only containing order-of-hundreds atoms. We have considered homogeneous ejecta for simplicity. In reality, ejecta may be inhomogeneous, and dust may be formed in dense clumps, as discussed in supernovae \citep[e.g.,][]{Kotak2009ApJ704p306,Indebetouw2014ApJ782L2}. Larger dust grains may be formed in higher density clumps, and then opacity by dust can be changed by reflecting the spatial distribution of dust. Moreover, the consideration of the radial profile of ejecta may modify a dust formation history. Such effects are interesting subjects to be studied in the future. The dust model for NIR macronovae should be tested observationally. One is the confirmation of a featureless spectrum. As shown in Figure \ref{fig:spec}, a dust emission spectrum is even featureless compared to a broad spectrum in the r-process model. A spectrum in the dust model also deviates from a blackbody spectrum at long wavelengths. Another way is multi-wavelength observations of light curves from early epochs (see Figure \ref{fig:lc}). Without opacity of r-process elements, a macronova is bright and blue in an early phase. It becomes red later by the emission of newly formed dust. Early optical emission was explored for a few short GRBs \citep[e.g., GRB 050509B;][]{Hjorth2005ApJ630L117}. Although the flux limit is strong for these GRBs, continuous searches for early emission are important because the properties of ejecta in NSB mergers may not be universal, e.g., depending on progenitors (NS-NS / BH-NS). In both cases, quick follow-up observations are important to understand the origin of NIR macronovae as well as the nucleosynthesis in the mergers of compact stellar objects.	14	3	1403.5872	We examine dust formation in macronovae (as known as kilonovae), which are the bright ejecta of neutron star binary mergers and one of the leading sites of r-process nucleosynthesis. In light of information about the first macronova candidate associated with GRB 130603B, we find that dust grains of r-process elements have difficulty forming because of the low number density of the r-process atoms, while carbon or elements lighter than iron can condense into dust if they are abundant. Dust grains absorb emission from ejecta with an opacity even greater than that of the r-process elements, and re-emit photons at infrared wavelengths. Such dust emission can potentially account for macronovae without r-process nucleosynthesis as an alternative model. This dust scenario predicts a spectrum with fewer features than the r-process model and day-scale optical-to-ultraviolet emission.	false	[ "Such dust emission", "emission", "r-process elements", "r-process nucleosynthesis", "infrared wavelengths", "elements", "Dust grains", "dust grains", "neutron star binary mergers", "dust formation", "ultraviolet", "dust", "neutron star binary", "the r-process model", "the r-process elements", "ejecta", "the r-process atoms", "GRB 130603B", "an alternative model", "macronovae" ]	6.430712	2.129221	62
437688	[ "Appourchaux, T.", "Antia, H. M.", "Benomar, O.", "Campante, T. L.", "Davies, G. R.", "Handberg, R.", "Howe, R.", "Régulo, C.", "Belkacem, K.", "Houdek, G.", "García, R. A.", "Chaplin, W. J." ]	2014A&A...566A..20A	[ "Oscillation mode linewidths and heights of 23 main-sequence stars observed by Kepler" ]	44	[ "Univ. Paris-Sud, Institut d'Astrophysique Spatiale, UMR 8617, CNRS, Bâtiment 121, 91405, Orsay Cedex, France", "Tata Institute of Fundamental Research, Homi Bhabha Road, 400005, Mumbai, India", "Sydney Institute for Astronomy (SIfA), School of Physics, University of Sydney, New South Wales, 2006, Sydney, Australia; Department of Astronomy, The University of Tokyo, 113-033, Tokyo, Japan", "School of Physics and Astronomy, University of Birmingham, Edgbaston, Birmingham, B15 2TT, UK; Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, 8000, Aarhus C, Denmark", "School of Physics and Astronomy, University of Birmingham, Edgbaston, Birmingham, B15 2TT, UK; Laboratoire AIM, CEA/DSM-CNRS-Université Paris Diderot, IRFU/SAp, Centre de Saclay, 91191, Gif-sur-Yvette Cedex, France", "School of Physics and Astronomy, University of Birmingham, Edgbaston, Birmingham, B15 2TT, UK; Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, 8000, Aarhus C, Denmark", "School of Physics and Astronomy, University of Birmingham, Edgbaston, Birmingham, B15 2TT, UK", "Instituto de Astrofísica de Canarias, 38205 La Laguna, Tenerife, Spain; Universidad de La Laguna, Dpto. de Astrofísica, 38206 La Laguna, Tenerife, Spain", "LESIA, Observatoire de Paris, CNRS UMR 8109, UPMC, Université Denis Diderot, 5 place Jules Janssen, 92195, Meudon Cedex, France", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, 8000, Aarhus C, Denmark", "Laboratoire AIM, CEA/DSM-CNRS-Université Paris Diderot, IRFU/SAp, Centre de Saclay, 91191, Gif-sur-Yvette Cedex, France", "School of Physics and Astronomy, University of Birmingham, Edgbaston, Birmingham, B15 2TT, UK" ]	[ "2014A&A...571A..71C", "2014A&A...572A..11G", "2015A&A...579A..83C", "2015EPJWC.10101008K", "2015LRSP...12....8H", "2015MNRAS.452.2654B", "2016A&A...589A.103R", "2016A&A...592A..14R", "2016A&A...595C...2A", "2016MNRAS.456.2183D", "2016MNRAS.462.1577Y", "2017A&A...601A..82W", "2017ApJ...835..172L", "2017ApJ...847...97S", "2017MNRAS.472..979H", "2017RAA....17...44L", "2018A&A...616A..94V", "2018A&A...617A...2M", "2018ApJ...857..119B", "2018ApJS..237...15A", "2018ApJS..239...34B", "2018MNRAS.476..470L", "2018MNRAS.478...69A", "2018arXiv180600994Y", "2019A&A...622A..76M", "2019A&A...623A.125B", "2019A&A...624A.117S", "2019A&A...624A.140S", "2019LRSP...16....4G", "2019MNRAS.482.1231J", "2019MNRAS.487..595H", "2019MNRAS.488..572K", "2020A&A...640A.130C", "2020A&A...641A..25N", "2020MNRAS.495.2363L", "2020MNRAS.495.4904Z", "2021A&A...650A.115D", "2021AJ....161...62N", "2022A&A...663A..51N", "2022A&A...663A.118B", "2022ApJ...928..188D", "2022FrASS...968452S", "2023A&A...680A..27B", "2024FrASS..1156379S" ]	[ "astronomy" ]	9	[ "stars: interiors", "asteroseismology", "methods: data analysis", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1980LNP...125..273G", "1980PASP...92..385V", "1982ApJ...263..835S", "1985ESASP.235..199H", "1988ApJ...334..510L", "1990ApJ...364..699A", "1992MNRAS.255..603B", "1994A&A...289..649T", "1994ApJ...424..480G", "1995A&A...293...87K", "1995SoPh..162..101F", "1996PhDT........80H", "1997IAUS..181...67F", "1997MNRAS.288..623C", "1997SoPh..170...27A", "1998ESASP.418..237H", "1999A&A...346..111L", "1999A&A...351..582H", "1999ApJ...513L..49F", "2000ApJ...543..472K", "2001MNRAS.327..483H", "2003ApJ...589.1009G", "2005A&A...433..713T", "2006MNRAS.369.1281B", "2007A&A...475..717M", "2008A&A...488..705A", "2008SoPh..251..189C", "2009A&A...500L..21C", "2009A&A...506...15B", "2009ApJ...694..144F", "2010ApJ...713L..87J", "2010ApJ...713L.160G", "2011A&A...527A..56H", "2011A&A...528A..25S", "2011A&A...529A..84B", "2011A&A...530A.142B", "2011A&A...531A.124B", "2011ApJ...731...78T", "2011LNP...832..305S", "2011MNRAS.414L...6G", "2011Natur.471..608B", "2011Sci...332..205B", "2011SoPh..274...87M", "2012A&A...537A.134A", "2012A&A...540A.143M", "2012A&A...540L...7B", "2012A&A...543A..54A", "2012A&A...548A..10M", "2012ApJ...756...19D", "2012ApJ...757..190C", "2012ApJS..199...30P", "2012MNRAS.423..122B", "2012Natur.481...55B", "2013ApJ...767..158B", "2013MNRAS.433.2893P", "2013MNRAS.433.3227K", "2014ApJS..210....1C", "2014aste.book..123A" ]	[ "10.1051/0004-6361/201323317", "10.48550/arXiv.1403.7046" ]	1403	1403.7046_arXiv.txt	\label{sec:intro} Stellar physics is undergoing a revolution thanks to the wealth of asteroseismic data that have been made available by space missions such as CoRoT \citep{Baglin2006} and \textit{Kepler} \citep{Gilliland2010}. With the seismic analyses of these stars providing the frequencies of the stellar eigenmodes, asteroseismology is rapidly becoming a tool for understanding stellar physics. Solar-type stars have been observed over periods exceeding six months using CoRoT and \textit{Kepler} providing many lists of mode frequencies required for seismic analysis \citep[See][and references therein]{Appourchaux2012a}. Additional invaluable information about the evolution of stars is provided by the study of the internal structure of red giants \citep{Bedding2011a, Beck2011, Beck2012, Mosser2012,Mosser2012a} and of sub giants \citep{Deheuvels2012, Benomar2013}. The large asteroseismic database of \textit{Kepler} allowed us to estimate the properties of an ensemble of solar-type stars that is large enough to perform statistical studies \citep{Chaplin2013a}. Solar-like oscillations are stochastically excited and damped by the convection. Thus measurements of mode linewidths and mode heights provide information about how the stellar modes are excited and damped. The processes involved are related to the generation of the acoustic noise and the dissipation of energy at the surface of the star \citep[See][]{GH99,Samadi2011}. For solar-like stars, a scaling relation for mode linewidth related to the stellar effective temperature has been proposed by \citet{Chaplin2009} using ground-based observations, \citet{Baudin2011} using CoRoT data and Appourchaux et al. (2012a) using \textit{Kepler} data. Those relations are based upon the linewidth measured at the frequency of maximum mode height and have been found by \citet{Belkacem2012} to be in qualitatively good agreement with the theoretical predictions. The relation was extended to lower effective temperature for red giants \citep{Corsaro2012}. These previous scaling studies do not provide the frequency dependence of the linewidth. Using a simple modelling approach, \citet{Gough1980} suggested that solar linewidths might have a local decrease, or dip, at the frequency of maximum power. With more accurate and detailed modelling, \citet{Balmforth1992} showed that there is indeed such a linewidth dip for the Sun. \citet{GH99} found that stellar mode linewidths show either a dip or {\it plateau} close to the maximum of mode height. The plateau is located at the frequency of the maximum of the mode height as shown by \citet{Belkacem2011}, which is also related to the Mach number (${\cal M}_a$), the ratio of convective velocity to the sound speed. This dip was first observed but not acknowledged in the solar p-mode linewidths by \citet{Libbrecht1988}, while a small dip or plateau was indeed observed by \citet{Chaplin1997}. The dip is caused by a resonance between the thermal adjustment time of the superadiabatic boundary layer and the mode frequency \citep{Balmforth1992}. Since the thermal adjustment time is proportional to the acoustic cut-off frequency $\nu_{c}$, the two frequencies follow a scaling relation as shown by \citet{Belkacem2011}. \citet{Claus1997} observed a very pronounced dip during solar minimum, hypothesizing that the depth of the dip may be modulated by solar activity, % as confirmed later by \citet{Komm2000}. These variations of the solar mode damping with solar activity were thought to be related to the change of the solar granule properties with the increasing magnetic field \citep{Houdek2001, Muller2007}, but these changes were not confirmed using space-based data \citep{Muller2011}. The variations are likely to be affected by the change in the global magnetic field during a solar cycle. Very recently, \citet{Benomar2013} studied the frequency dependence of mode linewidth of 4 sub-giant stars having mixed modes. They found that the linewidth of $l=0$ modes showed a clear dip at the location of the maximum of mode power. It was shown by \citet{Appourchaux2012} that different fits of the same data could provide significantly different results for stellar linewidths. Understanding the source of systematic errors will result in a better understanding of how physics operate in stars. \citet{Appourchaux2012} provided some insight on the various sources of systematic errors related to stellar background estimation and the mode height ratio. \citet{Chaplin2008a} showed that biased linewidths are also obtained when measuring mode linewidth of the order of 1 to 7 times larger than the frequency resolution. Apart from these two papers, the understanding of the origin of systematic errors on mode linewidth and height has not been widely studied. This paper aims at providing the frequency dependence of mode linewidth $\Gamma$, mode height $H$ and mode amplitude $A$ ($A=\sqrt{\pi H \Gamma /2}$) for 23 \textit{Kepler} main-sequence observed for nearly 2 years by \textit{Kepler}, as well as an understanding of the source of systematic errors affecting these parameters. The paper also aims at providing the dependence of the linewidth dip as a function of effective temperature. Section 2 describes how the time series and power spectra were obtained. Section 3 describes the peak fitting procedure. Section 4 details the sources of systematic errors on the mode linewidth and mode heights provided by the fitters. Section 5 provides a procedure for correcting the systematic errors with respect to a reference fit. We then discuss the detection of the dip as a function of effective temperature and the implications for stellar physics. The paper includes two examples of mode linewidth and mode height and an example of systematic error correction, while tables of the parameters of the 23 stars and correction for 22 stars are available online.	We have analysed the oscillation power spectra of 23 main-sequence stars for which we obtained the mode linewidths, mode heights and mode amplitude parameters. The parameters were obtained by a team of 8 independent fitters. We found large systematic errors between the parameters that could be traced to the way that the stellar background of the power spectrum was modelled; and to the fitted values of the rotational splitting and the stellar inclination angle. Other sources of systematic errors related to the mean frequency definition and to the mode height ratio were also studied. Finally using a correction scheme derived from the one-fit approach of \citet{TT2005b}, we could explain all sources of systematic errors, which could be reduced to less than $\pm$15\% for mode linewidth and mode height, and to less than $\pm$5\% for amplitude, when compared to a reference fit value. A different stellar background will give rise to frequency-dependent systematic errors that might affect the comparison with theoretical mode linewidth and mode height, therefore affecting the understanding of the physical nature of these parameters. All other sources of relative systematic errors are independent of frequency. Using the 23 stars of this study, 4 additional sub-giant stars of \citet{Benomar2013} and solar data, we also derived that the amplitude of the linewidth dip close to the maximum of frequency decreases with effective temperature. The dependence of the dip with effective temperature is linked to the behaviour of convection in the stellar atmosphere, implying that either the mixing length or the level of activity may increase with effective temperature. \begin{figure}[htbp] \centering \hbox{ \includegraphics[width=7 cm,angle=90]{linewidth_fit_1.ps} \includegraphics[width=7 cm,angle=90]{linewidth_fit_2.ps} } \caption{Parameters of Eq.~(\ref{Eq_fit}) as a function of the effective temperature for all 28 stars, for the power law dependence (Left) and for the Lorentzian fit (Right). The median value together with the credible intervals at 33\% and 66\% were derived from a Monte-Carlo simulation of the fit. The orange line shows the temperature dependence of the linewidth at the frequency of maximum mode height as derived by \citet{Appourchaux2012}. The open diamond is the result of the fit for the solar data of the LOI. The open triangles are the results of the fit for the sub-giant stars of \citet{Benomar2013}. The solid lines show a linear fit of the parameters with respect to the effective temperature. The Lorentzian parameters for stars for which the Lorentzian fit is not significant are not plotted.} \label{dip} \end{figure} \begin{figure}[htbp] \centering \hbox{ \includegraphics[width=7 cm,angle=90]{linewidth_fit_1yes.ps} \includegraphics[width=7 cm,angle=90]{linewidth_fit_2yes.ps} } \caption{Parameters of Eq.~(\ref{Eq_fit}) as a function of the frequency of maximum mode height for all 28 stars, for the power law dependence (Left) and for the Lorentzian fit (Right). The median value together with the credible intervals of 33\% and 66\% were derived from a Monte-Carlo simulation of the fit. The open diamond is the result of the fit for the solar LOI data. The open triangles are the result of the fit for the sub-giant stars of \citet{Benomar2013}. The error bars are derived from a Monte-Carlo simulation using credible intervals of 33\% and 66\%. The solid lines show a linear fit of the parameters with respect to the frequency. The Lorentzian parameters for stars for which the Lorentzian fit is not significant are not plotted.} \label{dip1} \end{figure}	14	3	1403.7046	Context. Solar-like oscillations have been observed by Kepler and CoRoT in many solar-type stars, thereby providing a way to probe the stars using asteroseismology. <BR /> Aims: We provide the mode linewidths and mode heights of the oscillations of various stars as a function of frequency and of effective temperature. <BR /> Methods: We used a time series of nearly two years of data for each star. The 23 stars observed belong to the simple or F-like category. The power spectra of the 23 main-sequence stars were analysed using both maximum likelihood estimators and Bayesian estimators, providing individual mode characteristics such as frequencies, linewidths, and mode heights. We study the source of systematic errors in the mode linewidths and mode heights, and we present a way to correct these errors with respect to a common reference fit. <BR /> Results: Using the correction, we can explain all sources of systematic errors, which could be reduced to less than ±15% for mode linewidths and heights, and less than ±5% for amplitude, when compared to the reference fit. The effect of a different estimated stellar background and a different estimated splitting will provide frequency-dependent systematic errors that might affect the comparison with theoretical mode linewidth and mode height, therefore affecting the understanding of the physical nature of these parameters. All other sources of relative systematic errors are less dependent upon frequency. We also provide the dependence of the so-called linewidth dip in the middle of the observed frequency range as a function of effective temperature. We show that the depth of the dip decreases with increasing effective temperature. The dependence of the dip on effective temperature may imply that the mixing length parameter α or the convective flux may increase with effective temperature. <P />Tables 4-27 and Appendices are available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201323317/olm">http://www.aanda.org</A>	false	[ "increasing effective temperature", "effective temperature", "mode linewidths", "mode height", "mode heights", "theoretical mode linewidth", "individual mode characteristics", "α", "various stars", "the mixing length parameter", "systematic errors", "relative systematic errors", "linewidths", "heights", "frequencies", "frequency", "Bayesian estimators", "the convective flux", "the mode linewidths", "many solar-type stars" ]	7.523953	11.776564	81
595893	[ "Deligny, Olivier" ]	2014CRPhy..15..367D	[ "Cosmic rays around 10<SUP>18</SUP> eV: Implications of contemporary measurements on the origin of the ankle feature" ]	12	[ "IPN Orsay, 15, rue Clémenceau, 91406 Orsay cedex, France" ]	[ "2014JCAP...11..031A", "2014JCAP...11..031B", "2014JPhCS.531a2009A", "2014arXiv1409.6138Z", "2015PhRvD..92l3001U", "2017PhRvD..96j3003A", "2018arXiv181103062A", "2019APh...104...13D", "2019JCAP...03..011A", "2020PhRvL.125l1106A", "2022ApJ...936...62L", "2023JCAP...05..024A" ]	[ "astronomy", "physics" ]	5	[ "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1934PNAS...20..259B", "1966PhRvL..16..748G", "1967PhLA...24..677H", "1970PhRvD...1.1596B", "1978JPhG....4..133E", "1978MNRAS.182..147B", "1978MNRAS.182..443B", "1987PhR...154....1B", "1991JPhG...17..733L", "1992JPhG...18..423N", "1993A&A...268..726P", "1993PhRvL..71.3401B", "2002APh....17..125S", "2003JCAP...05..003C", "2003NIMPA.513..490A", "2004APh....20..641A", "2004APh....21..583A", "2004APh....21..617B", "2004NuPhS.136..139H", "2005APh....24....1A", "2005PhRvD..71h3007L", "2005PhRvD..72h1301D", "2006PhRvD..74d3005B", "2007APh....27...61A", "2007ApJ...661L.175B", "2008JCAP...10..033A", "2010APh....33..151H", "2010PhLB..685..239A", "2010PhRvL.104i1101A", "2010PhRvL.104p1101A", "2011APh....34..627P", "2011ApJ...742..114P", "2011PhRvL.107q1104A", "2012APh....36...31B", "2012ApJ...760...48N", "2012ApJ...760..148P", "2012ApJS..203...34P", "2012JCAP...01..010B", "2012JCAP...01..011B", "2012JCAP...07..031G", "2012NIMPA.692...98B", "2013ApJ...762L..13P", "2013ApJ...768L...1A", "2013EPJWC..5301006B", "2013PhRvD..87h1101A", "2014ApJ...781...47K", "2014JCAP...10..020A" ]	[ "10.1016/j.crhy.2014.02.009", "10.48550/arXiv.1403.5569" ]	1403	1403.5569_arXiv.txt	\label{sec:intro} The ankle is a hardening of the energy spectrum of cosmic rays in the $10^{18}~$eV energy range. Discovered in 1963 by Linsley at the Volcano Ranch experiment~\cite{Linsley1963}, subsequent large-aperture experiments have nevertheless been necessary to characterise accurately this feature ~\cite{Lawrence1991,Nagano1992,Bird1993,AugerPLB2010,TA2013}. The ankle is today commonly described as a \textit{sharp} slope change of the spectral index of the cosmic ray intensity from $\simeq 3.3$ to $\simeq 2.7$ located at $\simeq 4\times 10^{18}~$eV. Although the existence of the ankle is beyond controversy, its interpretation is still under debate. A dedicated picture explaining this feature was already put forward by Linsley while he was reporting on this finding~: \og There are many possible interpretations. The one we favor is that the inflection marks a crossover between galactic and metagalactic cosmic rays\fg~\cite{Linsley1963}. In other words, this is nothing else but saying that the ankle may be the spectral feature marking the \textit{transition between Galactic and extragalactic cosmic rays}. The reasons supporting this interpretation are the object of section~\ref{sec:endgalactic}. Despite the fact that it has been popular for many years, this interpretation appears now in tension with contemporary measurements related to the mass composition of cosmic rays, the large-scale distribution of their arrival directions, and the energy spectrum for different mass group elements available around $10^{17}~$eV. The descriptions of these tensions are also the object of section~\ref{sec:endgalactic}. Alternatively, the ankle may be understood as the natural distortion of a proton-dominated extragalactic spectrum due to $e^{\pm}$ pair production in the collisions with the photons of the cosmic microwave background~\cite{Hillas1967,Blumenthal1970,Berezinsky2006,Berezinsky2004}. This requires on the one hand that the transition between Galactic and extragalactic cosmic rays takes place at lower energies and produces another spectral feature below $10^{18}~$eV~\cite{Berezinsky2006,Berezinsky2004}, and on the other hand that the mass composition is made of protons exclusively from the transition energy up to the highest ones. This scenario, referred to as \textit{the dip model}, is presented in section~\ref{sec:dip}. The Telescope Array and the Pierre Auger Observatory, which are currently running, are the two largest aperture experiments ever built to study cosmic rays around $10^{18}~$eV and above. Although the dip model provides an overall satisfactory theoretical framework in interpreting data collected at the Telescope Array, a totally different picture is needed in interpreting the ones collected at the Pierre Auger Observatory. This picture, though not final, is the object of section~\ref{sec:augerdata}. Hence, establishing the energy at which the intensity of cosmic rays starts to dominate the intensity of Galactic ones is still, as of today, a fundamental question in astroparticle physics. Part of the answer lies in understanding the ankle, but the whole picture will clearly appear only by revealing the origin of cosmic rays down to $10^{17}~$eV. To this end, some signatures of different scenarios that would be interesting, though challenging, to address in the future from an observational point of view are discussed in the final section.	\begin{figure}[!t] \centering \includegraphics[width=10cm]{phases_ironknee-ankle.eps} \caption{Phase of the first harmonic modulation in right ascension as a function of energy, from~\cite{Sidelnik2013,Curcio2013}.} \label{fig:phases} \end{figure} Despite the fact that the ankle was discovered in the early sixties, the origin of the ankle in the energy spectrum of cosmic rays is still a hot topic. Given the radically different interpretations of this feature which are possible from the data collected at the two contemporary large-aperture observatories - namely the Telescope Array and the Pierre Auger Observatory, uncovering and understanding the sources of systematic uncertainties in the respective measurements of the energy spectrum and of the atmospheric depths where the showers reach the maximum of their development is of central importance. Large-scale anisotropies are another part of the story which might help us understanding what is going on. Although no significant amplitude has been detected so far between $10^{17}$ and $10^{18}~$eV, it is to be noted that the phase measurements in adjacent energy intervals do not seem to be randomly distributed but rather to follow a constant behaviour as seen in figure~\ref{fig:phases}. This is potentially indicative of a genuine signal, because with a real underlying anisotropy, a consistency of the phase measurements in ordered energy intervals is indeed expected to be revealed with a smaller number of events than needed to detect the amplitude with high statistical significance~\cite{Edge1978,AugerAPP2011}. Whether this consistency provides a real evidence for anisotropy is currently being tested at the Pierre Auger Observatory, and will be established with 99\% confidence level if the ongoing prescribed test is successful~\cite{AugerApJS2012,Sidelnik2013}. \begin{figure}[!t] \centering \includegraphics[width=10cm]{toyanis_sig_vs_bkg.eps} \caption{Illustration of a mechanism authorising a large first-harmonic amplitude for the end of the Galactic component - see text.} \label{fig:toyanis} \end{figure} The putative dipole-like signal responsible for the phase alignment is required to have an amplitude rather low, being below the percent level to meet the conditions fixed by the upper limits obtained from the whole population of cosmic rays detected between $10^{17}$ and $10^{18}~$eV. Although hardly predictable in a quantitative way given the unknown source distributions in space and time and the difficulty to model the propagation of particles around $10^{18}~$eV in the intervening magnetic fields known only approximately, such an amplitude seems too low to be naturally explained by a Galactic scenario. On the other hand, isotropy down to a higher level is a robust expectation in this energy range for an extragalactic scenario. A positive outcome of the prescribed test would thus bring new pieces of information to the global puzzle. Different possibilities are discussed below. An interesting possibility to explain such a low global level of anisotropy could be that the component of heavy elements marking the end of the Galactic cosmic rays is slowly extinguishing from $10^{17}$ up to $\simeq 10^{18}~$eV, and does show important large-scale anisotropies in its distribution of arrival directions. Accounting for both the distribution of supernova remnants in the Galaxy and the random distribution in space and time of the nearest remnants, and considering the propagation of cosmic rays in a purely turbulent magnetic field, recent results \textit{extrapolated} above $10^{17}~$eV suggest as a sensible possibility an amplitude of the equatorial component of the dipole at the level of a few percents~\cite{Amato2012}. At these energies, the expected amplitudes could be amplified by drift motions induced by the regular component of the magnetic field~\cite{Ptuskin1993,Candia2003}. The anisotropy of this subdominant component of iron nuclei could be diluted in the almost isotropic predominant component of light elements - possibly extragalactic protons - so that the anisotropy of the all-particle component could satisfy the upper limits obtained between $10^{17}~$eV and $10^{18}~$eV. To illustrate this scenario, let's consider the total flux $\Phi_{\mathrm{tot}}(E,\mathbf{n})$ as the sum of a dominant isotropic component $\Phi_X(E)$ and of a subdominant anisotropic component $\Phi_{\mathrm{Fe}}(E,\mathbf{n})$ of iron-nuclei elements~: \begin{equation} \Phi_{\mathrm{tot}}(E,\mathbf{n})=\Phi_X^0E^{-\gamma_X}+\Phi_{\mathrm{Fe}}^0E^{-\gamma_{\mathrm{Fe}}}(1+r_{\mathrm{Fe}}(E)~\mathbf{d}\cdot\mathbf{n}). \end{equation} The anisotropy of the iron component is here characterised by an energy-dependent amplitude $r_{\mathrm{Fe}}(E)$ and a direction $\mathbf{d}$. The anisotropy amplitude of the total flux is then diluted as follows: \begin{equation} r_{\mathrm{tot}}(E)=\frac{\Phi_{\mathrm{Fe}}^0E^{-\gamma_{\mathrm{Fe}}}}{\Phi_X^0E^{-\gamma_X}+\Phi_{\mathrm{Fe}}^0E^{-\gamma_{\mathrm{Fe}}}}r_{\mathrm{Fe}}(E). \end{equation} For reasonable choices of the normalisation parameters and spectral indexes such that $\gamma_{\mathrm{Fe}}$ is larger than $\gamma_{\mathrm{X}}$, a \textit{decreasing} amplitude $r_{\mathrm{tot}}(E)$ with energy can be naturally obtained in this scenario even for an anisotropy of the iron component increasing with energy. This is illustrated in figure~\ref{fig:toyanis} for two \textit{toy} energy evolutions of the anisotropy inspired from diffusion-dominated ($r_{\mathrm{Fe}} \propto E^{0.3}$) or drift-dominated ($r_{\mathrm{Fe}} \propto E$) propagation of cosmic rays in the Galactic magnetic field. Although lots of fine-tunings are involved in this simplistic illustration, the global picture depicted in this plot may not be too irrealistic. Another possibility to mention is that the anisotropy of the iron component is not diluted, but almost \textit{compensated} by an anisotropy of the same order imprinted in the arrival directions of the component of light elements. Since the anisotropies discussed here are mainly described in terms of dipoles, this condition can be met by considering that the vectors characterising the dipoles of each component have an amplitude of the same order but an almost \textit{opposite} direction. This scenario provides a mechanism to reduce significantly the amplitude of the vector describing the arrival directions of the whole population of cosmic rays - which is observationally the only one at reach so far. This could relax to some extent the constraints on a Galactic origin of the new component. Along these lines, it is interesting to note the change of phase observed around $10^{18}~$eV in figure~\ref{fig:phases} between $\simeq 260^\circ$ and $\simeq 80^\circ$ in right ascension. Alternatively, if the iron-nuclei component is strongly cut off right after $10^{17}~$eV, the putative anisotropy would pertain to the component containing light elements. An anisotropy at the percent level in the $10^{18}~$eV energy range would challenge an extragalactic origin of this new component for equal-intensity sources spread over the Universe. Few strong local sources would be necessary to produce a strong density gradient of cosmic rays in the neighborhood of the Galaxy, so that an anisotropy could be observed on the condition that the pattern at the entrance of the Galaxy is not totally washed out by the Galactic magnetic field. On the other hand, a Galactic origin of this component was already shown to be in tension with the larger anisotropy amplitudes expected from stationary sources densely distributed predominantly in the disk and emitting in all directions - unless the strength of the Galactic magnetic field is much higher than currently known and/or much more extended perpendicularly to the disk. It is to be noted, however, that approximate anisotropy estimates from \textit{strongly intermittent} Galactic sources (\textit{e.g.} long gamma-ray bursts) show that the upper limits might be respected~\cite{Eichler2011,Kumar2013}. In this scenario, the observation of this extra-component of Galactic origin would be indicative of the remnant of an old and mostly isotropised population of cosmic rays from strongly intermittent sources \textit{different} from the ones producing the bulk of Galactic cosmic rays at lower energies. Large fluctuations in the intensity of the energy spectrum and in the anisotropies around $10^{18}~$eV are then intrinsic to this scenario. Current observations would not be representative of the usual ones and would be possible only during rare time periods. This could have some anthropic implications, which are however out of the scope of this report. That the ankle feature has not yet revealed all its secrets is indubitable. The end of the story will come from a comprehensive understanding of the origin of the different components observed between the iron knee and the ankle energies. An important step forward from an observational point of view would be to get at energy spectrum and anisotropy measurements for different mass group elements. This requires investing many experimental efforts to identify the primary mass on an event-by-event basis. A fair amount of exciting work remains.	14	3	1403.5569	The impressive power-law decay of the energy spectrum of cosmic rays over more than thirty orders of magnitude in intensity and for energies ranging over eleven decades between ≃10<SUP>9</SUP> eV and ≃10<SUP>20</SUP> eV is actually dotted with small irregularities. These irregularities are highly valuable for uncovering and understanding the modes of production and propagation of cosmic rays. They manifest themselves through changes in the spectral index characterising the observed power laws. One of these irregularities, known as the ankle, is subject to conflicting interpretations for many years. If contemporary observations characterising it have shed new lights, they are still far from being able to deliver all the story. The purpose of this contribution is to give an overview of the physics of cosmic rays in the energy range where the transition between Galactic and extragalactic cosmic rays is expected to occur, and to deliver several lines of thought about the origin of the ankle.	false	[ "cosmic rays", "extragalactic cosmic rays", "many years", "conflicting interpretations", "small irregularities", "energies", "several lines", "≃10", "thought", "intensity", "magnitude", "eV", "Galactic", "new lights", "propagation", "the energy range", "production", "the energy spectrum", "uncovering", "the ankle" ]	6.186451	0.421579	13
526255	[ "Hamada, Yuta", "Kawai, Hikaru", "Oda, Kin-ya", "Park, Seong Chan" ]	2014arXiv1403.5043H	[ "Higgs inflation still alive" ]	52	[ "-", "-", "-", "-" ]	[ "2014EPJC...74.3022L", "2014GReGr..46.1786N", "2014IJMPA..2950099H", "2014JCAP...06..032R", "2014JCAP...06..039F", "2014JHEP...06..010B", "2014PhLB..734...41G", "2014PhLB..734..249B", "2014PhLB..735..186K", "2014PhRvD..89j3525C", "2014PhRvD..89k5009H", "2014SCPMA..57.1460C", "2014arXiv1403.6403O", "2014arXiv1404.1538O", "2014arXiv1405.2743K", "2014arXiv1405.2931D", "2014arXiv1408.7079G", "2015EPJC...75...55L", "2015EPJC...75..301L", "2015EPJC...75..355Z", "2015JHEP...01..128I", "2015PTEP.2015e3B01H", "2015PTEP.2015i1B01H", "2015PTEP.2015l3B03H", "2015arXiv150406707K", "2015arXiv150406931I", "2016EPJC...76..303C", "2017FoPh...47..769A", "2018FrASS...5...40R", "2018JCAP...06..005E", "2018arXiv180706830A", "2018arXiv180709952P", "2018arXiv181000792S", "2018arXiv181010546G", "2019FrASS...6...55R", "2019OJAp....2E...1R", "2019arXiv190513581D", "2020JHEP...05..154M", "2020Physi...2..503H", "2020arXiv200808700H", "2020arXiv201007563M", "2021Univ....7..354S", "2021arXiv211101362L", "2022EPJC...82..504S", "2022JHEP...02..100L", "2022JHEP...06..164I", "2022JHEP...09..101B", "2022arXiv221211977J", "2023JHEP...10..068R", "2023JHEP...10..144C", "2023arXiv231011260C", "2023arXiv231201718H" ]	[ "astronomy", "physics" ]	11	[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]	[ "1992eaqg.book.....B", "1996PhLB..368...96F", "2001PhRvD..64k3014F", "2006PhRvL..97s1304A", "2007PhLB..648..312M", "2008JCAP...08..009P", "2008PhLB..659..703B", "2008PhLB..660..260M", "2008PhRvD..77b5034I", "2008PhRvD..77c5006F", "2009JHEP...07..089B", "2009JHEP...09..103B", "2009PhLB..676...81I", "2009PhRvD..79h1302B", "2009PhRvD..80k5007I", "2010JHEP...07..007B", "2010PhLB..683..196S", "2011JHEP...01..016B", "2011PhLB..694..294G", "2011PhRvD..83h3515K", "2012JHEP...02..037H", "2012JHEP...08..098D", "2012JHEP...10..140B", "2012PThPh.127..689K", "2012PhLB..716..214A", "2012PhRvD..86a3001A", "2012PhRvD..86b3504K", "2012arXiv1212.5716N", "2013CQGra..30u4001B", "2013JHEP...12..089B", "2013PTEP.2013b3B08I", "2013PhLB..727..234S", "2013PhRvD..87e3001M", "2013PhRvD..87e3009H", "2013PhRvD..87f4021G", "2013PhRvD..88e6022B", "2013PhRvD..88i3003M", "2013PhRvD..88l3518K", "2013arXiv1305.7055H", "2013arXiv1309.7335J", "2013arXiv1310.0563H", "2014AcPPB..45.1167J", "2014JCAP...02..024G", "2014JHEP...02..040A", "2014JHEP...04..029H", "2014JHEP...06..010B", "2014PDU.....5...75M", "2014PTEP.2014b3B02H", "2014PTEP.2014f3B01I", "2014PhLB..734...96N", "2014PhRvD..89a6019H", "2014PhRvD..89e6010H", "2014PhRvD..89j3525C", "2014PhRvL.112a1302K", "2014PhRvL.112x1101B", "2014arXiv1403.4427A" ]	[ "10.48550/arXiv.1403.5043" ]	1403	1403.5043_arXiv.txt			14	3	1403.5043	The observed value of the Higgs mass indicates that the Higgs potential becomes small and flat at the scale around $10^{17}$GeV. Having this fact in mind, we reconsider the Higgs inflation scenario proposed by Bezrukov and Shaposhnikov. It turns out that the non-minimal coupling $\xi$ of the Higgs-squared to the Ricci scalar can be smaller than ten. For example, $\xi=7$ corresponds to the tensor-to-scalar ratio $r\simeq0.2$, which is consistent with the recent observation by BICEP2.	false	[ "Higgs", "Shaposhnikov", "BICEP2", "mind", "Bezrukov", "the Higgs inflation scenario", "scalar", "the Higgs potential", "Ricci", "the Higgs mass", "the Ricci scalar", "the recent observation", "this fact", "the scale", "example", "$r\\simeq0.2", "around $10^{17}$GeV.", "$\\xi$", "The observed value", "$r\\simeq0.2$" ]	10.607619	-1.192784	89
745644	[ "Umehata, H.", "Tamura, Y.", "Kohno, K.", "Hatsukade, B.", "Scott, K. S.", "Kubo, M.", "Yamada, T.", "Ivison, R. J.", "Cybulski, R.", "Aretxaga, I.", "Austermann, J.", "Hughes, D. H.", "Ezawa, H.", "Hayashino, T.", "Ikarashi, S.", "Iono, D.", "Kawabe, R.", "Matsuda, Y.", "Matsuo, H.", "Nakanishi, K.", "Oshima, T.", "Perera, T.", "Takata, T.", "Wilson, G. W.", "Yun, M. S." ]	2014MNRAS.440.3462U	[ "AzTEC/ASTE 1.1-mm survey of SSA22: Counterpart identification and photometric redshift survey of submillimetre galaxies" ]	51	[ "Institute of Astronomy, The University of Tokyo, Mitaka, Tokyo 181-0015, Japan", "Institute of Astronomy, The University of Tokyo, Mitaka, Tokyo 181-0015, Japan", "Institute of Astronomy, The University of Tokyo, Mitaka, Tokyo 181-0015, Japan; Research Center for the Early Universe (WPI), University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "North American ALMA Science Center, National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, Virginia 22903, USA", "Astronomical Institute, Tohoku University, 6-3 Aoba, Aramaki, Aoba-ku, Sendai, Miyagi 980-8578, Japan", "Astronomical Institute, Tohoku University, 6-3 Aoba, Aramaki, Aoba-ku, Sendai, Miyagi 980-8578, Japan", "UK Astronomy Technology Centre, Science and Technology Facilities Council, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ; Institute for Astronomy, University of Edinburgh, Blackford Hill, Edinburgh EH9 3HJ", "Department of Astronomy, University of Massachusetts, Amherst, MA 01003, USA", "Instituto Nacional de Astrofisica, Optica y Electronica (INAOE), Aptdo. Postal 51 y 216, 72000 Puebla, Pue., Mexico", "Center for Astrophysics and Space Astronomy, University of Colorado, Boulder, CO 80309, USA", "Instituto Nacional de Astrofisica, Optica y Electronica (INAOE), Aptdo. Postal 51 y 216, 72000 Puebla, Pue., Mexico", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "Astronomical Institute, Tohoku University, 6-3 Aoba, Aramaki, Aoba-ku, Sendai, Miyagi 980-8578, Japan", "Institute of Astronomy, The University of Tokyo, Mitaka, Tokyo 181-0015, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan; Nobeyama Radio Observatory, National Astronomical Observatory of Japan, Minaminaki, Minamisaku, Nagano 384-1305, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan; Joint ALMA Observatory, Alonso de Cordova 3107, Vitacura, Santiago 763 0355, Chile", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "Advanced Technology Center, National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan; Joint ALMA Observatory, Alonso de Cordova 3107, Vitacura, Santiago 763 0355, Chile; The Graduate University for Advanced Studies (Sokendai), Mitaka, Tokyo 181-8588, Japan", "Nobeyama Radio Observatory, National Astronomical Observatory of Japan, Minaminaki, Minamisaku, Nagano 384-1305, Japan", "Department of Astronomy, University of Massachusetts, Amherst, MA 01003, USA", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "Department of Astronomy, University of Massachusetts, Amherst, MA 01003, USA", "Department of Astronomy, University of Massachusetts, Amherst, MA 01003, USA" ]	[ "2014ApJ...797..138C", "2014MNRAS.443..146J", "2015A&A...577A..29M", "2015ASPC..499...29U", "2015ApJ...799...38K", "2015ApJ...815L...8U", "2015MNRAS.448.3325J", "2016A&ARv..24...14O", "2016ApJ...820...82C", "2016ApJ...826..130H", "2016IAUS..319..109U", "2016MNRAS.455.2363H", "2016MNRAS.455.3333K", "2016MNRAS.460.3861K", "2016MNRAS.461.2944A", "2016PASJ...68...82Y", "2017ApJ...835...98U", "2017ApJ...835..286I", "2017ApJ...850..178A", "2017FrASS...4...51J", "2017MNRAS.467.3951M", "2017MNRAS.469.2235K", "2018ApJ...856...72O", "2018PASJ...70...65U", "2019ApJ...871...83S", "2019ApJ...883..142H", "2019ApJ...887..214K", "2019MNRAS.487.4648S", "2020A&A...640L...8U", "2020IAUS..352..157U", "2020MNRAS.495.2332K", "2020MNRAS.498.3095Y", "2020PASJ...72...69Y", "2021ApJ...907..122M", "2021ApJ...919....6K", "2021ApJ...919...51M", "2021MNRAS.500..942D", "2021MNRAS.501.1803L", "2021arXiv211011977M", "2022A&A...660A.137G", "2022A&A...664A.155G", "2022ApJ...929..159C", "2022ApJ...930...32C", "2022MNRAS.512.4893Z", "2022Univ....8..554A", "2023AJ....165..208M", "2023ApJ...951...15M", "2023ApJ...953...75H", "2023ApJ...958...36S", "2024ApJ...960..132L", "2024MNRAS.528.2964Z" ]	[ "astronomy" ]	6	[ "galaxies: distances and redshifts", "galaxies: formation", "galaxies: high-redshift", "large-scale structure of Universe", "submillimetre: galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]	[ "1980ApJ...236..351D", "1986MNRAS.218...31D", "1992ARA&A..30..575C", "1993ApJ...405..538B", "1993ApJ...412...64L", "1996A&AS..117..393B", "1997ApJ...490L...5S", "1998ApJ...492..428S", "1998Natur.394..241H", "1998Natur.394..248B", "1999ApJ...515..518E", "1999MNRAS.303..659H", "2000A&A...363..476B", "2000ApJ...532..170S", "2000ApJ...533..682C", "2002PASJ...54..833M", "2004AJ....128.2073H", "2004ApJ...606...85C", "2004ApJ...611..725B", "2004ApJS..154...10F", "2004ApJS..154...25R", "2004MNRAS.354..779G", "2004SPIE.5489..763E", "2005ApJ...622..772C", "2005ApJ...623..721P", "2005ApJ...634L.125M", "2005MNRAS.359.1165G", "2005Natur.435..629S", "2005PASP..117.1113M", "2006MNRAS.372.1621C", "2007MNRAS.377.1717K", "2007MNRAS.379.1571A", "2007MNRAS.379.1599L", "2007MNRAS.380..199I", "2008MNRAS.385.2225S", "2008MNRAS.386..807W", "2008MNRAS.389..333Y", "2008PASJ...60..683U", "2008PASJ...60.1347S", "2008SPIE.7012E..08E", "2009ApJ...690.1236I", "2009ApJ...691..560C", "2009ApJ...694.1517D", "2009ApJ...699.1610H", "2009MNRAS.395.1905C", "2009MNRAS.400..299L", "2009Natur.459...61T", "2010Ap&SS.330..219S", "2010ApJ...721.1056S", "2010ApJ...724.1270T", "2011ApJ...733...92W", "2011MNRAS.411..102H", "2011MNRAS.413.2314B", "2011MNRAS.415.1479W", "2011Natur.470..233C", "2012A&A...548A...4S", "2012AJ....143...79Y", "2012ApJ...750..116U", "2012ApJ...761...89B", "2012MNRAS.420..957Y", "2012MNRAS.421..284H", "2012MNRAS.423..529D", "2012MNRAS.423..575S", "2012MNRAS.424.2232A", "2012MNRAS.426.1845M", "2012MNRAS.427.2866S", "2013ApJ...768...91H", "2013ApJ...776..131C", "2013ApJ...778..170K", "2013MNRAS.430.2768T", "2013MNRAS.431..194A", "2013MNRAS.432....2K" ]	[ "10.1093/mnras/stu447", "10.48550/arXiv.1403.2725" ]	1403	1403.2725_arXiv.txt	The first deep extragalactic survey at 850 $\mu$m undertaken with the Submillimeter Common-User Bolometer Array (SCUBA, \citealt{1999MNRAS.303..659H}) on the James Clerk Maxwell Telescope (JCMT) unveiled a population of galaxies in the distant redshift universe that are extremely bright at submillimeter wavelengths (\citealt{1997ApJ...490L...5S}, \citealt{1998Natur.394..241H}, \citealt{1998Natur.394..248B}). This was followed by several wide surveys at (sub)millimeter wavelengths have been conducted to discover more and more such galaxies (e.g., \citealt{2004MNRAS.354..779G}, \citealt{2008MNRAS.385.2225S}, \citealt{2009MNRAS.395.1905C}, \citealt{2011MNRAS.411..102H}). These submillimeter bright galaxies (SMGs) have huge rest-frame far-infrared (FIR) luminosities ($L_{\mathrm{FIR}}\sim10^{12}-10^{13}L_{\odot}$), which should be caused mainly by highly dust-enshrouded star formation and are often indicative of a star-formation rate (SFR) of $\geq$ 1000 M$_\odot$/yr. Their flux at submillimeter wavelengths is almost constant for galaxies with a given FIR luminosity at $z\sim1$--10 due to the negative $k$-correction. It is hence of great benefit to discover high redshift objects (for review, see \citealt{2002PhR...369..111B}). In addition to their high activity, spectroscopic observations of the millimeter and submillimeter transitions of molecular carbon monoxide (CO) have unveiled the SMGs' large dynamical and gas masses (e.g., \citealt{2005MNRAS.359.1165G}). This observational evidence shows that SMGs were the most active, massive star forming galaxies in the early universe. The clarification of the SMG nature and formation process is seriously important to understand the galaxy formation and evolution. As with other galaxy populations, one of the most crucial questions about the SMGs is how their formation and evolution depend on their environment. The current cold dark matter (CDM) cosmological simulations show that the SMGs should preferentially exist in regions where the mass densities are high and, correspondingly, also the merger rates are high (e.g., \citealt{2005Natur.435..629S}). SMGs are also supposed to be progenitors of the massive elliptical galaxies seen in the cores of present-day rich clusters (e.g., \citealt{1999ApJ...515..518E}). While the connection between SMGs and massive dark matter haloes has been statistically indicated by clustering analysis (e.g., \citealt{2004ApJ...611..725B}; \citealt{2012MNRAS.421..284H}), the connection in individual cluster/protocluster is still unclear, although previous papers have reported some cases. \cite{2011Natur.470..233C} and \cite{2009ApJ...694.1517D} reported the discovery of SMGs in overdense region at $z=5.3$ and $z=4.05$, respectively. These results indicate that the overdense regions might be sites of SMG formation. On the other hand, \cite{2009ApJ...691..560C} shows that SMGs are formed in less overdence regions at $z=1.99$. Thus the environmental dependence on SMG formation is still controversial. The relationship between SMGs and the surrounding environment is also intriguing from the view point of galaxy formation at (proto)clusters. In the local universe the morphology-density relation has been well-known. Observations have revealed a higher fraction of early-type galaxies in denser environments \citep{1980ApJ...236..351D}. This trend has been confirmed for up to $z\sim$1 universe (e.g., \citealt{2005ApJ...623..721P} ). Although this is one of the most established environmental effects on galaxy evolution, it is difficult to examine the relation directly at higher redshifts. Instead, the color-density or color-magnitude relations were examined as proxies. For instance, \cite{2007MNRAS.377.1717K} examined the color-magnitude relation in protoclusters and found that the red sequence of galaxies, which is well-established in clusters at least out to $z\sim1$, first appeared at $z=2-3$. This suggests that massive galaxies were assembled in protoclusters in this era. SSA22 is a unique laboratory field to investigate the formation of star-forming galaxies including SMGs in the overdense region. \cite{1998ApJ...492..428S} first discovered this protocluster as a concentration of Lyman break galaxies (LBGs) at $z = 3.09$. Furthermore \cite{2000ApJ...532..170S} found that the surface density of Lyman alpha emitters (LAEs) was also much higher than that of other fields. Consequently, the wide field survey using Subaru/Suprime-Cam equipped with a narrowband filter (NB497) have revealed a structure which was traced by LAEs and spread over 700 arcmin$^2$(\citealt{2004AJ....128.2073H}; \citealt{2012AJ....143...79Y}). \cite{2012AJ....143...79Y} evaluated that the degree of overdensity was at most 10 times of the expected standard deviations based on the counts of LAEs at $z=3.1$. Hence the SSA22 field can provide us with unique insights regarding galaxy formation in an overdense environment. Several previous works on SMGs discovered by SCUBA and AzTEC/ASTE surveys in this field has been reported. \cite{2004ApJ...611..725B} and \cite{2005ApJ...622..772C} confirmed three SCUBA SMGs have $z_{spec}=3.1$ and these SMGs really lie within the densest region. \cite{2009Natur.459...61T} showed that there was an angular correlation between the 15 brightest AzTEC SMGs and $z=3.1$ LAEs. But the lack of redshift information prevents us from investigating this further. Although redshift is one of the essential pieces of information for this purpose, obtaining them has remained a difficult task. Firstly, the typical beam size of the single-dish telescopes used for wide field SMG surveys is insufficient to determine position accurately. In the case of the AzTEC/ASTE survey, we can achieve only $\sim30^{\prime\prime}$ FWHM. Accurate positions of SMGs could ideally be obtained with a sub-millimeter interferometer, like the Atacama Large Millimeter/submillimeter Array (ALMA; \citealt{2010SPIE.7733E..34H}), but such observations are time-intensive over large fields. Secondly SMGs generally tend to be optically faint due to dust attenuation and hence optical observations are often helpless to determine counterpart and obtain redshift information. To overcome such difficulties in searching counterparts and determining redshift, previous works have shown that multi-wavelength identification utilizing radio, MIPS(\citealt{2004ApJS..154...25R}), and IRAC(\citealt{2004ApJS..154...10F}) imaging data is useful (e.g., {\citealt{2011MNRAS.413.2314B}; \citealt{2011MNRAS.415.1479W}; \citealt{2012MNRAS.420..957Y}). If the identified counterparts in these images have an optical to near-infrared counterpart, we can derive a photometric redshift. This is the approach we will follow. We expand the area concerned by Tamura et al. and add photometric redshift information to investigate the relation between SMGs and underlying environment more closely. The organization of this paper is as follows. In Section 2 we report the general properties of the AzTEC/ASTE survey and the detected SMGs. The utilization of other wavelength data set of this field is presented in Section 3. In Section 4 we describe our analysis used for counterpart identification. The estimation of the photometric redshifts are derived in Section 5. In Section 6 we discuss the relationship between SMGs and the $z=3.1$ protocluster. We assume a cosmology with $\Omega_m=0.3, \Omega_\lambda=0.7, H_0=70$ km s$^{-1}$ Mpc$^{-1}$ and all magnitudes are given according to the AB system throughout the paper.	We imaged a 950 arcmin$^2$ field towards the SSA22 field covering a 1$\sigma$ depth down to $0.7-1.3$ mJy using the AzTEC 1.1mm camera attached on the ASTE. This survey area corresponds to about 2.5 times of the previously reported area (390 arcmin$^2$) by \cite{2009Natur.459...61T}. We detected 125 SMGs with S/N $\ge$ 3.5 and eight out of them are expected to be fake source arising from noise peaks. We attempted to identify reliable counterparts to 125 AzTEC SMGs utilizing VLA 1.4 GHz imaging, MIPS 24 $\mu$m images and IRAC color diagnostics (\citealt{2008MNRAS.389..333Y}). We considered the corrected Poissonian probability, $p$, to evaluate the degree of chance coincidence. We regarded counterpart candidates as robust counterparts if $p$ is less than 0.05. Additionally sources with $0.05\leq p<0.20$ were considered as tentative counterparts. We found that 59 SMGs have at least one reliable (i.e., robust or tentative) counterpart considering all diagnostics methods. We performed SED fitting based on optical to near-infrared photometry utilizing the {\it HYPERZ} code, which provided us with photometric redshifts for 48 counterparts of 45 SMGs, which were all covered by all band IRAC observations. We couldn't find any counterparts for 14 of 61 SMGs that had four-band IRAC coverage. The redshift distribution of the SSA22 field is similar to those of GOODS-S AzTEC SMGs reported by \cite{2012MNRAS.420..957Y}. These AzTEC SMGs tend to lie at a higher redshift universe than the radio-identified SCUBA SMGs ($z_{med}=2.2$ reported by \citealt{2005ApJ...622..772C}). Some AzTEC sources without reliable counterparts may be located at higher-$z$ and would enhance this trend. We found 10 AzTEC SMGs that possibly are at $z=3.1$ based on photometric and/or spectroscopic redshifts. Among them, seven out of the 10 SMGs were concentrated into the core 12 Mpc $\times$ 12 Mpc region (comoving scale), which is consistent with the center of the galaxy distribution of the protocluster. Cross-correlation functions indicate that the distribution of these 10 SMGs and that of $z=3.1$ LAEs are correlated. These results suggest that SMGs are tend to be formed in extremely high density environments.	14	3	1403.2725	We present the results from a 1.1-mm imaging survey of the SSA22 field, known for having an overdensity of z = 3.1 Lyman α emitting galaxies (LAEs), taken with the astronomical thermal emission camera (AzTEC) on the Atacama Submillimeter Telescope Experiment (ASTE). We imaged a 950-arcmin<SUP>2</SUP> field down to a 1σ sensitivity of 0.7-1.3 mJy beam<SUP>-1</SUP> to find 125 submillimetre galaxies (SMGs) with a signal-to-noise ratio ≥3.5. Counterpart identification using radio and near/mid-infrared data was performed and one or more counterpart candidates were found for 59 SMGs. Photometric redshifts based on optical to near-infrared images were evaluated for 45 of these SMGs with Spitzer/IRAC data and the median value is found to be z = 2.4. By combining these estimations with estimates from the literature, we determined that 10 SMGs might lie within the large-scale structure at z = 3.1. The two-point angular cross-correlation function between LAEs and SMGs indicates that the positions of the SMGs are correlated with the z = 3.1 protocluster. These results suggest that the SMGs were formed and evolved selectively in the high dense environment of the high-redshift Universe. This picture is consistent with the predictions of the standard model of hierarchical structure formation.	false	[ "z", "ASTE", "hierarchical structure formation", "mid-infrared data", "SMGs", "the Atacama Submillimeter Telescope Experiment", "the astronomical thermal emission camera", "AzTEC", "LAEs", "Lyman", "noise", "Photometric redshifts", "α emitting galaxies", "Universe", "Spitzer", "125 submillimetre galaxies", "the high dense environment", "Spitzer/IRAC data", "10 SMGs", "59 SMGs" ]	13.493918	7.110414	132
483083	[ "Lee, Jong Hwan", "Lee, Myung Gyoon" ]	2014ApJ...786..130L	[ "A New Optical Survey of Supernova Remnant Candidates in M31" ]	50	[ "Astronomy Program, Department of Physics and Astronomy, Seoul National University, Seoul 151-747, Korea", "Astronomy Program, Department of Physics and Astronomy, Seoul National University, Seoul 151-747, Korea" ]	[ "2014ApJ...793..134L", "2014SerAJ.189...15G", "2015ApJ...804...63L", "2015MNRAS.446..943V", "2017ApJ...844...69M", "2017ApJS..230....2B", "2017MNRAS.464.2326S", "2017MNRAS.472..308G", "2017arXiv170107840L", "2017hsn..book.2005L", "2017hsn..book.2087B", "2017suex.book.....B", "2018A&A...620A..28S", "2018ApJ...855..140L", "2018ApJ...861...92D", "2018ApJS..239...13W", "2018MNRAS.473.4130M", "2018MNRAS.480.3052M", "2018MNRAS.481.2804E", "2018RNAAS...2...32D", "2018arXiv180902034N", "2019ApJ...877...15G", "2019ApJ...887...66P", "2019ApJS..241...37W", "2019MNRAS.483.4551E", "2020A&A...634A...3V", "2020AN....341..156S", "2020MNRAS.491..889K", "2020MNRAS.492..848A", "2020MNRAS.498.5367D", "2021A&A...651A..98F", "2021AJ....162..199L", "2021ApJ...908...80W", "2021ApJ...908...85M", "2021ApJ...920...62N", "2021MNRAS.505.5301R", "2021MNRAS.507.6020K", "2021NewA...8301492E", "2021NewA...8301498R", "2022AJ....163...60W", "2022ApJ...928...54S", "2023A&A...672A.148C", "2023AJ....165..116L", "2023ApJ...949...32K", "2023ApJ...954..107C", "2023ApJ...959...62W", "2023ApJS..265...53R", "2023ApJS..268...36H", "2024MNRAS.530.1078K", "2024arXiv240508974L" ]	[ "astronomy" ]	16	[ "galaxies: individual: M31", "galaxies: ISM", "ISM: supernova remnants", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1958ApJ...128..465D", "1964IAUS...20..245M", "1972ARA&A..10..129W", "1977ApJ...218..148M", "1978A&A....63...63D", "1979ApJS...39....1R", "1980A&AS...40...67D", "1981ApJ...247..879B", "1983ApJS...51..345M", "1984ApJ...276..653D", "1984AuJPh..37..321M", "1984BAAS...16..927H", "1984MNRAS.209..449G", "1988AJ.....96.1874C", "1990ApJS...72...61L", "1991ApJ...375..652S", "1993A&AS...98..327B", "1995A&AS..114..215M", "1995AJ....110..739L", "1996ApJ...466..750L", "1997ApJS..108..261B", "1997ApJS..112...49M", "1997ApJS..113..333M", "1998A&AS..130..421F", "1998ApJ...500..342B", "1998ApJS..117...89G", "1999AJ....118.2775G", "1999ApJS..120..299T", "1999ApJS..123..467W", "2000ASPC..221...83S", "2002ApJ...565..966P", "2003adu..book.....D", "2004A&A...421.1031V", "2004A&A...426...11P", "2004ApJS..155..101B", "2005A&A...434..483P", "2005AJ....130..539G", "2005MNRAS.364..217F", "2006A&A...448.1247M", "2006AJ....131.2478M", "2006ApJ...638L..87G", "2007AJ....133.1361P", "2007ApJ...663..234G", "2008ApJS..174..366P", "2008MNRAS.383.1175P", "2009A&A...493.1061S", "2009ApJ...700..727B", "2009BASI...37...45G", "2010A&A...509A..34B", "2010A&A...509A..70V", "2010ApJ...710..964D", "2010ApJS..187..495L", "2010MNRAS.407.1301B", "2011A&A...534A..55S", "2011piim.book.....D", "2012A&A...544A.144S", "2012AJ....143...85F", "2012ApJ...745..156R", "2012ApJ...761...26J", "2012ApJS..203....8B", "2013ApJS..204....4P", "2013MNRAS.429..189L" ]	[ "10.1088/0004-637X/786/2/130", "10.48550/arXiv.1403.4335" ]	1403	1403.4335_arXiv.txt	Supernova remnants (SNRs) play an important role in our understanding of supernovae (SNe), the interstellar medium (ISM), and the interaction between them. Large samples of SNRs in a galaxy can be used to understand the evolution of SNRs, estimate the SN rate in galaxies, and investigate the global properties of the ISM in the galaxy as well as the local ISM. SNRs are generally divided into two categories according to their progenitors: core-collapse (CC) and Type Ia SNRs. CC SNRs result from CC SNe caused by massive stars undergoing core collapse, while Type Ia SNRs are remnants of Type Ia SNe occurring when a white dwarf (WD) accretes material from its binary companion, causing the WD mass to exceed the Chandrasekhar limit. These two types of SNe eject different mixtures of heavy elements into the ISM of a galaxy, which have different impacts on the galactic chemical evolution. Therefore, a study of the properties of these two types of SNRs in galaxies can provide a clue to understand the star formation history and chemical evolution of galaxies. There are 274 known SNR candidates in our Galaxy; thus, our Galaxy has the largest sample of known SNR candidates in the universe \citep{gre09}. However, they occupy too large an angular size in the sky and their distances are not well known, so it is difficult to obtain their optical properties. Therefore, the information on the statistical properties of these SNR candidates is very limited. On the other hand, SNR candidates in nearby galaxies do not suffer from these problems, so they are an ideal target for studying the optical properties statistically. Extragalactic SNR surveys have been conducted at optical, radio, and X-ray wavelengths. The first SNR candidates were identified from a radio survey of the Large Magellanic Cloud (LMC) \citep{mat64}. Since then, 77 SNR candidates have been identified in the MCs using radio, X-ray, and optical techniques \citep{fil98,wil99,smi00,van04,pay08,bad10}. However, limitations on the sensitivity and resolution reduce the effectiveness of radio and X-ray searches for SNRs in nearby galaxies. Therefore, optical searches have produced the largest number of extragalactic SNRs. Previous optical surveys identified $\sim$230 SNR candidates in M31, $\sim$140 SNR candidates in M33, and several hundred SNR candidates in other nearby galaxies using photometric and spectroscopic data \citep{mag95,mat97a,mat97b,pan02,bla04,son09,lon10,dop10,fra12}. Recently, \citet{bla12} found 225 SNR candidates in M83 using \ha and \s2 images obtained at the Magellan I 6.5 m telescope, and \citet{leo13} identified $\sim$ 400 SNR candidates in five nearby galaxies using narrow-band images obtained at the 1.4 m telescope. However, SNR surveys in galaxies beyond the Local Group are limited by the available sensitivity and resolution. Because SNRs in these galaxies are typically unresolved in ground-based images, their morphologies are largely unknown. Many of the previously known SNR candidates have sizes of $D >$ 100 pc, which is larger than typical SNRs. Therefore, previous surveys might have included spurious SNRs such as \h2 and superbubbles. Additionally, they missed many faint and diffuse SNRs. \citet{dop10} found 60 SNR candidates in M83 using high-resolution $HST$ images, but they covered only a fraction of the galaxy. We started a project to study SNRs in nearby galaxies using wide-field optical images. In this study, we selected M31, which is an appealing galaxy for studying SNRs owing to its proximity ($\sim$750 kpc, \citealt{vil10,rie12}). At the distance of M31, an SNR with $D \sim$ 20 pc has an angular size of $\sim$5$\arcsec$. Therefore, it is possible to distinguish many SNRs using ground-based images, characterize them in detail, and classify them considering their morphological structures. M31 has a significant number of optically identified SNRs. \citet{dod80} identified 19 SNR candidates on the basis of their \ratb, and \citet{bla81} confirmed 14 SNR candidates with enhanced \rat using spectroscopic data. \citet{bra93} found 52 SNR candidates using narrow-band images in \ha and \s2, but their survey was limited to portions of the northwestern half of M31. Most of the known SNR candidates are credited to \citet{mag95} who reported 179 SNR candidates. However, \citet{mag95} did not cover the entire region of M31. They could not find faint SNRs because of the short exposure times. Additionally, because they used narrow-band images obtained under poor seeing conditions ($>$2\arcsec), they could not resolve the SNRs well. For example, they could not distinguish SNRs located around the outside edges of giant \h2. They included objects having large sizes ($D >$ 100 pc), which might be superbubbles. We conducted a new SNR survey of M31 using the data provided by the Local Group Survey (LGS) \citep{mas06}. M31 was observed as part of the LGS program with the KPNO/Mayall 4 m telescope in \hab, \s2, and \o3 as well as other continuum bands. The survey covered the entire disk of M31. In this study, we present the results of the SNR survey over the entire disk region in M31. This paper is composed as follows. Section 2 describes the data and explains the methods used to identify SNR candidates, measure their sizes and fluxes, and classify them considering their progenitors and morphology. Section 3 provides a catalog of M31 SNR candidates and presents their spatial distributions, \ha and \s2 luminosity functions, size distributions, \rat distributions, and radial distributions. In Section 4, we compare the optical properties of the M31 SNR candidates with those in other nearby galaxies, probe correlations between the optical properties and X-ray luminosity of the M31 SNR candidates, and investigate the difference between the distributions of Type Ia and CC SNR candidates. Finally, a summary and conclusion is given in Section 5.	We found 76 new SNR candidates through a wide-field survey based on \ha and \s2 images of M31 in the LGS. In addition, we confirmed that 80 of the 234 SNR candidates in previous studies are SNR candidates according to our selection criteria. In our analysis, we investigated various properties of the 156 SNR candidates in the master catalog. The primary results are summarized as follows. \begin{enumerate} \item We attempted to classify the progenitor types of our SNR candidates according to the properties of the stellar and interstellar populations in and around each candidate. We found that 42 more likely result from Type Ia SNe, with the remainder more likely to be from CC SNe. The fraction of Type Ia SNR candidates in M31 ($\sim$23\%) is similar to that found in M101 \citep{fra12}. \item We classified SNR candidates considering their optical morphologies as well as their general environments. The numbers of A-type, B-type, and C-type SNR candidates are 54, 85, and 17, respectively. The numbers of shell-type and center-bright SNR candidates are 133 ($\sim85\%$) and 23 ($\sim15\%$), respectively. These fractions are comparable to those for the MW SNRs. \item Most of the CC SNR candidates are concentrated in the spiral arms, while the Type Ia SNR candidates are rather spread over the entire galaxy including the inner region. The radial distribution of the CC SNR candidates shows two distinct peaks at $R \sim$ 12 kpc and $R \sim$ 5 kpc, while that of the Type Ia SNR candidates is broad, showing no distinct peaks. Most of the SNR candidates at $R <$ 7 kpc are A-type SNR candidates, while there are more B-type SNR candidates in the outer region at $R >$ 7 kpc. This indicates that the ISM at $R <$ 7 kpc may be more uniform than that in other regions. \item Most of the Type Ia SNR candidates have fainter \ha and \s2 luminosities than the CC SNR candidates. The \ha luminosity function of all the SNR candidates is fitted by a double power law with a break at $L \sim 10^{36.6}$\ergs. The power indices for the bright and faint parts are $\alpha = -2.61 \pm 0.42$ and $\alpha = -1.26 \pm 0.17$, respectively. The \s2 luminosity function is fitted by a single power law with an index of $\alpha = -2.24 \pm 0.03$. \item Most of the SNR candidates in M31 have sizes of 20 pc $< D <$ 60 pc. The differential size distribution of all the SNR candidates shows a strong peak at $D\sim $ 48 pc with a broad wing. The differential size distribution of the CC SNR candidates shows a strong peak at $D\sim $ 48 pc, while that of the Type Ia SNR candidates is much broader with a weaker peak at $D\sim $ 40 pc. The differential size distribution of the A-type SNR candidates shows a mean value of $D\sim $ 35 pc, which is much smaller than that of the B-type SNR candidates, $D\sim $ 60 pc. \item The cumulative size distribution of all the SNR candidates with 17 pc $< D <$ 50 pc is well fitted by a power law with an index of $\alpha = 2.53 \pm 0.04$. This indicates that most of the M31 SNR candidates identified in this study appear to be in the Sedov--Talyor phase. The cumulative size distribution of the CC SNR candidates with 15 pc $< D <$ 55 pc is fitted by a power law with $\alpha = 2.30 \pm 0.04$, and that of the Type Ia SNR candidates with 25 pc $< D <$ 50 pc is fitted by a similar power law with $\alpha = 2.45 \pm 0.06$. The difference in the size ranges between the two types of candidates indicates that most of the CC SNR candidates may lie in a denser ambient ISM than the Type Ia SNR candidates. The cumulative size distribution of the A-type SNR candidates with 25 pc $< D <$ 45 pc is fitted by a power law with $\alpha = 2.15 \pm 0.09$, while that of the B-type SNR candidates with 35 pc $< D <$ 60 pc is fitted by a much steeper power law with $\alpha = 4.63 \pm 0.14$. This indicates that the B-type SNR candidates may evolve faster than the A-type SNR candidates. \item The \rat distribution of all the SNR candidates is bimodal, with peaks at \rat $\sim$ 0.4 and $\sim$ 0.9. The \rat distributions of the CC and Type Ia SNR candidates are similarly bimodal. The ratio of CC SNR candidates and Type Ia SNR candidates is higher for low \ratb than for high \ratb. The high \rat populations are mostly A1-type, A2-type, B1-type, and B4-type SNR candidates that have well-defined shell-like or compact shapes, while the low \rat populations are mostly A3-type, B2-type, and B3-type SNR candidates that have low surface brightness and smooth shapes. The B2-type SNR candidates are embedded in star-forming regions, and they have high \ha luminosity and low \ratb. \item We inspected the radial variation in the physical properties of the SNR candidates. In the inner region ($R <$ 10 kpc), the mean size of the SNRs increases from 40 pc to 60 pc as $R$ increases, while the mean values of their \ha and \s2 surface brightnesses decrease. The \rat of all the SNR candidates shows little variation with $R$, which is in contrast to the result given by \citet{gal99}. \item The \ha and \s2 luminosities of all the SNR candidates show weak or little linear correlation with their sizes. Those of the CC SNR candidates show linear correlations with their sizes, while those of the Type Ia SNR candidates show little correlation. The \ha and \s2 surface brightnesses of all the SNR candidates show linear correlations with their sizes. Those of the Type Ia SNR candidates show stronger linear correlations with their sizes than those of the CC SNR candidates. The \ha and \s2 luminosities of each morphological type show tight linear correlations with their sizes. The \ha and \s2 surface brightnesses of each morphological type show little linear correlation with their sizes. \item The cumulative size distribution of the M31 SNR candidates with 17 pc $< D <$ 50 pc is well fitted by a power law with an index of $\alpha = 2.53 \pm 0.04$. The cumulative size distribution of the M33 SNR candidates with 13 pc $< D <$ 33 pc identified in \citet{gor98} is well fitted by a power law with an index of $\alpha = 2.72 \pm 0.14$. The result is similar to that for the M31 SNR candidates, although the mean size of the M33 SNR candidates following a Sedov--Taylor phase is smaller than that of the M31 SNR candidates. This suggests two possibilities. First, the incompleteness of large SNR detection in our study is lower than that for M33. Second, most of the M31 SNR candidates may lie in a less dense ambient ISM than the M33 SNR candidates. \item A higher fraction of SNR candidates ($\sim$21\% and $\sim$80\% for A1-type and A2-type, respectively) with relatively high surface brightnesses, small sizes, and complete shapes are detected in X-rays. We inspected the correlations between the optical properties and X-ray luminosities of the 23 SNR candidates in M31 common to this study and \citet{sas12}. We found a good correlation between the optical and X-ray luminosities for the combined sample of A1-type and A2-type SNR candidates, and a better correlation between the optical surface brightness and X-ray luminosity for the A2-type SNR candidates. These results indicate that the ambient medium of the SNR candidates with X-ray counterparts has a locally uniform density. \item The radial distribution of the SNR candidates with X-ray counterparts shows a distinct peak near 12 kpc and a broad peak near 5 kpc. Most of the SNR candidates at $R < 6$ kpc are Type Ia SNR candidates with X-ray counterparts. The X-ray luminosities of most of the Type Ia SNR candidates are brighter than those of the CC SNR candidates. \end{enumerate} The authors thank Prof. Bon-Chul Koo for a fruitful discussion on SNRs and an anonymous referee for useful comments that helped improve the original manuscript significantly. This work was supported by a National Research Foundation of Korea (NRF) grant funded by the Korea Government (MSIP) (No. 2012R1A4A1028713).	14	3	1403.4335	We present a survey of optically emitting supernova remnants (SNRs) in M31 based on Hα and [S II] images in the Local Group Survey. Using these images, we select objects that have [S II]:Hα > 0.4 and circular shapes. We identify 156 SNR candidates, of which 76 are newly found objects. We classify these SNR candidates according to two criteria: the SNR progenitor type (Type Ia and core-collapse (CC) SNRs) and the morphological type. Type Ia and CC SNR candidates make up 23% and 77%, respectively, of the total sample. Most of the CC SNR candidates are concentrated in the spiral arms, while the Type Ia SNR candidates are rather distributed over the entire galaxy, including the inner region. The CC SNR candidates are brighter in Hα and [S II] than the Type Ia SNR candidates. We derive a cumulative size distribution of the SNR candidates, finding that the distribution of the candidates with 17 <D < 50 pc is fitted well by a power law with the power-law index α = 2.53 ± 0.04. This indicates that most of the SNR candidates identified in this study appear to be in the Sedov-Taylor phase. The [S II]:Hα distribution of the SNR candidates is bimodal, with peaks at [S II]:Hα ~ 0.4 and ~0.9. The properties of these SNR candidates vary little with the galactocentric distance. The Hα and [S II] surface brightnesses show a good correlation with the X-ray luminosity of the SNR candidates that are center-bright.	false	[ "CC SNR candidates", "Type Ia", "the Type Ia SNR candidates", "CC SNR", "SNR", "The CC SNR candidates", "the CC SNR candidates", "the SNR candidates", "156 SNR candidates", "these SNR candidates", "Hα", "SNRs", "the SNR progenitor type", "S II]:Hα &gt", "[S II", "the X-ray luminosity", "α", "CC", "the candidates", "objects" ]	4.025633	5.201791	-1
484399	[ "Choi, Jieun", "Conroy, Charlie", "Moustakas, John", "Graves, Genevieve J.", "Holden, Bradford P.", "Brodwin, Mark", "Brown, Michael J. I.", "van Dokkum, Pieter G." ]	2014ApJ...792...95C	[ "The Assembly Histories of Quiescent Galaxies since z = 0.7 from Absorption Line Spectroscopy" ]	137	[ "Department of Astronomy & Astrophysics, University of California, Santa Cruz, CA 95064, USA", "Department of Astronomy & Astrophysics, University of California, Santa Cruz, CA 95064, USA", "Department of Physics and Astronomy, Siena College, Loudonville, NY 12110, USA", "Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA", "UCO/Lick Observatories, University of California, Santa Cruz, CA 95064, USA", "Department of Physics and Astronomy, University of Missouri, Kansas City, MO 64110, USA", "School of Physics, Monash University, Clayton, Vic 3800, Australia", "Department of Astrophysical Sciences, Yale University, New Haven, CT 06520, USA" ]	[ "2015A&A...573A..78Q", "2015A&A...573A.110G", "2015A&A...582A..46W", "2015ApJ...798L...4M", "2015ApJ...799..125V", "2015ApJ...800...94S", "2015ApJ...803...77C", "2015ApJ...803...87S", "2015ApJ...804L...4M", "2015ApJ...806L..31M", "2015ApJ...807...11G", "2015ApJ...807...36C", "2015ApJ...808..161O", "2015ApJ...808L..32S", "2015ApJ...811L..12W", "2015IAUS..311..126G", "2015MNRAS.448.1430M", "2015MNRAS.451.1158P", "2015MNRAS.454.1332S", "2016A&A...585A..86H", "2016A&A...592A..19C", "2016ApJ...820..120B", "2016ApJ...822....1F", "2016ApJ...831..173F", "2016ApJ...832...79P", "2016ApJ...833....2G", "2016ApJS..223...29V", "2016MNRAS.456.2140M", "2016MNRAS.457.3743D", "2016MNRAS.461.1131M", "2016MNRAS.463..832W", "2016Natur.540..248K", "2016pas..conf..246S", "2017A&A...597A.107S", "2017A&A...599A..95G", "2017AJ....154..251J", "2017ApJ...836..120R", "2017ApJ...841...32Z", "2017ApJ...843..105S", "2017ApJ...847...18Z", "2017ApJ...847...20W", "2017ApJ...851...34B", "2017MNRAS.468.1747C", "2018A&A...616A.121S", "2018A&A...617A.113E", "2018ASPC..517..639K", "2018ApJ...854..139C", "2018ApJ...855...85W", "2018ApJ...855..142E", "2018ApJ...856...15L", "2018ApJ...856L...4M", "2018ApJ...857...22L", "2018ApJ...861...13C", "2018ApJ...863..191J", "2018ApJ...868...84M", "2018ApJ...869..117R", "2018ApJS..239...27S", "2018MNRAS.475.4148G", "2018MNRAS.480.4379C", "2018PhDT.......124M", "2019A&A...625A..94H", "2019A&A...630A.145T", "2019A&A...631A.156D", "2019A&A...631A.157D", "2019A&ARv..27....3M", "2019AJ....158....2B", "2019ApJ...870..133E", "2019ApJ...872..136C", "2019ApJ...873...63Z", "2019ApJ...874...17B", "2019ApJ...875...16C", "2019ApJ...876....3L", "2019ApJ...877...48C", "2019ApJ...877..140L", "2019ApJ...877..141M", "2019ApJ...878..158Z", "2019ApJ...879...45V", "2019ApJ...880L..31K", "2019ApJ...881...42J", "2019ApJ...885..100L", "2019MNRAS.486.1358F", "2019MNRAS.490..417C", "2020A&A...644L...7G", "2020AJ....159..186D", "2020ApJ...889...93V", "2020ApJ...898...82M", "2020ApJ...900...95V", "2020MNRAS.491.2822C", "2020MNRAS.491.5406T", "2020MNRAS.498.5317W", "2020arXiv200309357D", "2021A&A...649A..93B", "2021AJ....162..201W", "2021ApJ...906...43P", "2021ApJ...908...84V", "2021ApJ...917L...1B", "2021ApJ...920...63Z", "2021FrASS...8..157D", "2021MNRAS.500.3368S", "2021MNRAS.504.2190E", "2022ApJ...924...25M", "2022ApJ...924...32V", "2022ApJ...926..134T", "2022ApJ...927..164B", "2022ApJ...929..131C", "2022ApJ...934..177K", "2022LRR....25....6M", "2022MNRAS.511..341C", "2022MNRAS.512.1262H", "2022MNRAS.512.3828B", "2022MNRAS.517.2697L", "2023A&A...672A..87I", "2023A&A...677A..93D", "2023ApJ...948...79M", "2023ApJ...948..132Z", "2023ApJ...948..140B", "2023ApJS..266...13S", "2023MNRAS.519.1476A", "2023MNRAS.520.3027S", "2023MNRAS.520.5651C", "2023MNRAS.520.6091D", "2023MNRAS.521.5400H", "2023MNRAS.523.4251A", "2023MNRAS.525.4219B", "2023MNRAS.526.4004C", "2023SCPMA..6629511W", "2023arXiv230709501M", "2024ApJ...961..118K", "2024ApJ...966..234B", "2024ApJS..270...12W", "2024MNRAS.528.6010U", "2024MNRAS.529.3342G", "2024MNRAS.530.4550P", "2024MNRAS.531.3858R", "2024arXiv240412432S", "2024arXiv240501637D", "2024arXiv240603549J" ]	[ "astronomy" ]	10	[ "galaxies: abundances", "galaxies: elliptical and lenticular", "cD", "galaxies: evolution", "galaxies: stellar content", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1970SAOSR.309.....K", "1974MNRAS.169..229L", "1979ApJ...229.1046T", "1984ApJ...286..644N", "1984ApJ...287..586B", "1990MNRAS.245..217G", "1993sssp.book.....K", "1994ApJS...94..687W", "1996A&AS..117..393B", "1996MNRAS.281..985V", "1998ApJS..116....1T", "1999ASPC..193..258J", "1999MNRAS.302..537T", "2000AJ....120..165T", "2000AJ....120.1579Y", "2000MNRAS.317..965H", "2001ApJ...553...90V", "2001MNRAS.322..231K", "2001MNRAS.328.1039C", "2002AJ....124.1810S", "2002ApJ...579L..13S", "2003ApJ...585..694E", "2003ApJ...590L..91S", "2003MNRAS.339L..12C", "2003MNRAS.343..279T", "2003MNRAS.343..978S", "2003PASP..115..763C", "2003SPIE.4841..525A", "2003SPIE.4841.1657F", "2004A&A...428.1043L", "2004AJ....128.2826W", "2004ApJ...608..752B", "2004MSAIS...5...93S", "2004PASP..116..138C", "2005A&A...439..845L", "2005AJ....129.2562B", "2005ApJ...619L...1M", "2005ApJ...621..673T", "2005ApJ...623..666R", "2005ApJ...626..680D", "2005ApJ...631..145V", "2005ApJ...633..174T", "2005MNRAS.358..363C", "2005PASP..117.1411F", "2006A&A...453L..29C", "2006A&A...457..787S", "2006ARA&A..44..141R", "2006ApJ...639..816E", "2006ApJ...648..268B", "2006ApJ...650...18T", "2006ApJ...651..120B", "2006ApJ...651..791B", "2006ApJ...651L..93S", "2006ApJ...653..159K", "2006MNRAS.365...46O", "2007A&A...467..991B", "2007A&A...476..137A", "2007ApJ...654..858B", "2007ApJ...655...30V", "2007ApJ...658..710N", "2007ApJ...665..265F", "2007ApJ...671..243G", "2007ApJS..169...21C", "2007ApJS..171..146S", "2007ApJS..172..494S", "2007MNRAS.374..769S", "2007MNRAS.377..759S", "2007MNRAS.378.1507B", "2007MNRAS.380..585C", "2008A&A...482...21C", "2008A&A...487...89V", "2008ApJ...677L...5V", "2008ApJ...688...48V", "2008ApJS..177..446G", "2008MNRAS.386..715T", "2009A&A...499...47S", "2009ApJ...691.1862M", "2009ApJ...691L..95T", "2009ApJ...694..902L", "2009ApJ...695..101D", "2009ApJ...697.1290B", "2009ApJ...699..486C", "2009ApJ...699L.178N", "2009ApJ...701..428A", "2009ApJ...701.1765M", "2009ApJ...702..307S", "2009ApJ...703..785D", "2009ApJ...704L..34C", "2009ApJS..182..543A", "2009ApJS..185....1T", "2009MNRAS.395..160K", "2009MNRAS.395..608T", "2009MNRAS.398..119S", "2009MNRAS.398..898H", "2009MNRAS.398L..44W", "2009Natur.459..814P", "2010A&A...519A..55E", "2010A&A...523A..13P", "2010AJ....140.1868W", "2010ApJ...709..644I", "2010ApJ...709.1018V", "2010ApJ...712..833C", "2010ApJ...713..738W", "2010ApJ...714L..79C", "2010ApJ...718.1460F", "2010ApJ...719..844R", "2010ApJ...719.1715W", "2010ApJ...722..491Z", "2010ApJ...724..714H", "2010ApJS..190..233M", "2010MNRAS.401..852R", "2010MNRAS.401..933M", "2010MNRAS.402.2264B", "2010MNRAS.404.1775T", "2010MNRAS.404.2087B", "2010MNRAS.408..272S", "2011AJ....142...31B", "2011Ap&SS.331....1W", "2011ApJ...727...51N", "2011ApJ...736L...9V", "2011ApJ...739...24B", "2011ApJ...741....8C", "2011MNRAS.414..445S", "2011MNRAS.415..709S", "2011MNRAS.417..900D", "2011arXiv1112.3300M", "2012A&A...548A...7L", "2012ApJ...747...69C", "2012ApJ...748...10C", "2012ApJ...753...44S", "2012ApJ...755...26O", "2012ApJS..200....8K", "2012MNRAS.421.1908J", "2012MNRAS.425..641L", "2012MNRAS.425..841L", "2012MNRAS.426.2300L", "2013AJ....145...77J", "2013ARA&A..51..393C", "2013ApJ...767...50M", "2013ApJ...773..112C", "2013ApJ...774...28B", "2013ApJ...776...64G", "2013ApJ...777...18M", "2013ApJ...779..138B", "2013ApJS..208....5N", "2013MNRAS.428..999P", "2013MNRAS.428.1088M", "2013MNRAS.429.2924H", "2013MNRAS.436..697B", "2013PASP..125..306F", "2014ApJ...780...33C", "2014ApJ...781...38C", "2014ApJ...783...20W", "2014ApJ...788...72G", "2014ApJ...788L..29B", "2014ApJ...794...65M" ]	[ "10.1088/0004-637X/792/2/95", "10.48550/arXiv.1403.4932" ]	1403	1403.4932_arXiv.txt	Considerable progress has been made over the last several decades toward understanding the formation and assembly histories of galaxies. On the observational front, two complementary techniques are widely used. Lookback studies aim to study the evolution of galaxies in a statistical manner by comparing snapshots of the galaxy population at different cosmic epochs. This technique has been powered by major spectroscopic surveys such as the Sloan Digital Sky Survey \citep[SDSS;][]{York2000}, the 2dF Galaxy Redshift Survey \citep[2dFGRS;][]{Colless2001}, the VIMOS VLT Deep Survey \citep[VVDS;][]{LeFevre2004, LeFevre2005}, the PRIsm MUlti-object Survey \citep[PRIMUS;][]{Coil2011}, the Deep Extragalactic Evolutionary Probe 2 (DEEP2) survey \citep{Newman2013}, and the AGN and Galaxy Evolution Survey \citep[AGES;][]{Kochanek2012}, among many others. In addition, targeted surveys out to $z\sim1$ have provided many insights into the assembly histories of galaxies \citep[e.g.,][]{vanDokkum1996, Rusin2005, Treu2005, Holden2010, Jorgensen2013}. In contrast, in the archaeological approach, one infers past evolution through detailed studies of $z\sim0$ galaxies and their stellar populations. This method of extrapolating back in time is enabled by high quality data of nearby galaxies. One of the most basic probes of galaxy formation and evolution enabled by large-scale redshift surveys is the time evolution of galaxy luminosity and stellar mass functions. At $z\gtrsim2$, star-forming galaxies dominate quiescent galaxies in number at all stellar masses, but the mass density of quiescent galaxies has increased by almost an order of magnitude between $z\sim2$ and today \citep[e.g.,][]{Bell2004, Blanton2006, Brown2007, Faber2007, Cirasuolo2007, Arnouts2007, Cappellari2009, Ilbert2010, Whitaker2010, vandeSande2011, Brammer2011, DominguezSanchez2011}. While the quiescent population has been on a global rise, there is strong evidence that this process is mass-dependent. Between $z\sim2$ and $z\sim1$, the space density of massive galaxies has been observed to increase rapidly \citep{Arnouts2007, Cirasuolo2007, Ilbert2010, Nicol2011, Brammer2011}. By $z\sim1$, massive galaxies are mostly assembled, and they appear to passively evolve to $z\sim0$ \citep[e.g.,][]{Bundy2006, Renzini2006, Cirasuolo2007, Vergani2008, Marchesini2009, Banerji2010, Moustakas2013, Muzzin2013}. Thus the rapid evolution in the mass and luminosity functions of quiescent galaxies at late times has mostly been attributed to the rise in low- and intermediate-mass quiescent galaxies, although the details are still under active debate \citep[e.g.,][]{Cimatti2006, Scarlata2007, Brown2007, Stewart2009, Ilbert2010, Robaina2010, Whitaker2010, ElicheMoral2010, Pozzetti2010, Brammer2011, Skelton2012, Moustakas2013}. However, mass and luminosity functions by necessity depend on the ability to accurately measure global photometry of extended sources. In particular, accurate photometry accounting for diffuse light at large radius is notoriously difficult even at $z=0$ \citep{Bernardi2013}, and it only becomes more challenging at higher redshifts due to the $(1+z)^{4}$ decrease in surface brightness. There has been a tremendous effort in the past decade to study the outskirts of massive quiescent galaxies as a function of cosmic time. Both size evolution studies using deep imaging and dynamical studies have shown that the sizes of massive galaxies have increased by a factor of $2\text{--}4$ from $z\sim2$ to the present \citep{Daddi2005, Trujillo2006, vanDokkum2008, vanderWel2008, Cimatti2008, Bezanson2009, Damjanov2009, Williams2010, Cassata2010, vanDokkum2010, LopezSanjuan2012, McLure2013, Belli2013}. The inner regions ($r\lesssim5$ kpc) of massive galaxies have apparently undergone very little mass growth, but mass out to $\sim75$ kpc has increased by a factor of four since $z\sim2$ \citep[e.g.,][]{vanDokkum2010}. The main channel for size and mass growth is likely dominated by minor mergers, though in-situ star formation and major mergers are also thought to play a role. A simple scaling argument based on the virial theorem explains the dramatic size growth during minor mergers, where the radius increases quadratically with mass instead of linearly as in major mergers \citep{Naab2009}. Observations showing that the oldest and most massive galaxies at high redshift are smaller by a factor of $\sim5$ compared to low-redshift galaxies of comparable mass \citep[e.g.,][]{Daddi2005, Onodera2012} further corroborate the ``inside-out growth'' of massive galaxies. Simulations support this notion that galaxy formation occurs in two phases: compact cores are thought to form rapidly at $z\gtrsim2$ from star formation triggered by infalling cold gas, followed by a slower growth in both mass and size over a longer period of time through the accretion of satellites \citep[e.g.,][]{Naab2007, Keres2009, Naab2009, Hopkins2009a, Dekel2009, Lackner2012, Hilz2013}. This is still a controversial field, however, with some groups proposing other modes of size growth, e.g., baryonic mass loss leading to the ``puffing up'' of galaxies \citep[e.g.,][]{Fan2010}. Others invoke progenitor bias, arguing that the new galaxies entering the quiescent population are inherently larger in size \citep[e.g.,][]{Carollo2013}, and some groups even find evidence for a lack of size evolution \citep[e.g.,][]{Mancini2010, Stott2011}. Stellar population analysis offers yet another way of probing the evolution of galaxies (see \citealt{Walcher2011}, \citealt{Conroy2013} for recent reviews). The most popular technique for analyzing properties of old stellar populations in quiescent galaxies, such as age, metallicity, and abundance patterns, is to measure and model several key absorption features using the Lick/IDS index system \citep{Burstein1984, Worthey1994}. Based on this approach, stellar populations in massive quiescent galaxies are found to be old and enhanced in $\alpha$ elements compared to the Milky Way disk stars \citep[e.g.,][]{Worthey1994}. Moreover, numerous independent groups have found strong positive correlations with velocity dispersion for age, total metallicity, and the ratio of $\alpha$ elements to Fe---almost ubiquitously represented by [Mg/Fe] in the literature---but almost no trend for [Fe/H] \citep{Trager1998, Thomas2005, Graves2007, Schiavon2007, Smith2009, Zhu2010, Johansson2012, Conroy2014, Worthey2014}. The total metallicity is an important diagnostic for galaxy formation and evolution because it is sensitive to the depth of the potential well in which their stellar populations formed \citep[e.g., supernovae-driven winds can efficiently remove metals from shallow potential wells;][]{Larson1974}. On the other hand, [$\alpha$/Fe] is sensitive to the timescale of star formation: massive stars expel $\alpha$ elements into their interstellar environments on million-year timescales, while Type Ia supernovae enrich the star-forming gas with Fe on billion-year timescales \citep[e.g.,][]{Tinsley1979}. By measuring the relative abundances of $\alpha$ elements and Fe in stellar populations, the time and the rate at which stars formed in their host galaxies can be inferred \citep[e.g.,][]{Thomas1999}. More recently, several groups have begun modeling the full optical spectrum of galaxies. First used to measure star formation histories and metallicities \citep{Heavens2000, CidFernandes2005, Ocvirk2006, Tojeiro2009}, this technique was further developed by \cite{Walcher2009} and \cite{Conroy2012} to include variable elemental abundances. The present work can be viewed as the high-redshift extension of \cite{Conroy2014}, which focused on full spectrum modeling of high-quality SDSS spectra of $z\sim0$ quiescent galaxies to derive their ages and detailed abundance patterns. In this work, we use a hybrid approach of combining lookback and archaeological studies to carry out detailed stellar population analysis at progressively higher redshifts. For the first time, we derive accurate ages and abundance measurements for a large mass-complete sample of galaxies from $z\sim0.1$ to $z\sim0.7$, and examine how their stellar population properties evolve over time as a function of stellar mass. As we demonstrate, powerful constraints can be placed on the assembly histories of galaxies from the time evolution of their stellar population properties. The paper is organized as follows: Section~\ref{section:data} gives an overview of the data sets and the sample selection process, and Section~\ref{section:model} presents background information on our models and the full spectrum fitting technique. The main science results are introduced in Section~\ref{section:results}, and a discussion and summary are provided in Sections~\ref{section:discussion} and \ref{section:summary}. We discuss the results of various systematic tests to explore the robustness of our analysis in the Appendix. All wavelengths in this paper are quoted in vacuum. Where necessary, we assume a Chabrier IMF \citep{Chabrier2003} for the stellar mass range \mbox{0.1--100~$\msun$} unless noted otherwise, and $\Lambda$CDM cosmology, with \mbox{$H_0$ = 70 \kms{} Mpc$^{-1}$}, $\Omega_{\rm m} = 0.3$, and $\Omega_{\rm \Lambda} = 0.7$. \begin{figure}[] \centering \includegraphics[width=1.7\columnwidth]{ages_sdss_ssfr_logM_sameplot-crop.pdf} \caption{{\it Left panel}: specific star formation rate (sSFR) as a function of stellar mass for a subset of the SDSS sample. A random subset is shown for display purposes. Both the quiescent cloud and the star-forming main sequence can be seen. We locate the minimum of the bimodal distribution in three mass slices and perform a linear least-squares fit to those points. The resulting quiescence threshold is shown as a dashed line. We discard all galaxies above this quiescence cut to obtain a sample of quiescent galaxies. {\it Right panel}: sSFR as a function of stellar mass for the entire AGES sample. The black dashed line corresponds to the quiescence cut determined from the SDSS sample. Only those that fall below the sSFR threshold are included in our quiescent sample. We use a quiescence cut that is constant with redshift because there is no evidence for evolution in the location of the green valley.} \label{fig:sSFR_vs_mass} \end{figure}	\label{section:discussion} \subsection{The Assembly Histories of Quiescent Galaxies} Our results indicate that for quiescent galaxies, there is negligible evolution in elemental abundances ($\lesssim0.1$~dex) at fixed stellar mass over roughly 7~Gyr, and the increase in their ages with cosmic time is consistent with passive evolution since $z=0.7$ (see Figures~\ref{fig:ages_sdss_mass} and \ref{fig:logage_redshift}). The stellar masses from SED-fitting nominally represent the total stellar mass. However, since our stellar population analysis is based on spectra obtained with a $1\farcs5$ fiber, the measured parameters are sensitive to stellar populations in the inner regions of galaxies directly probed by the fiber. There are two effects at play: the angular size of galaxies of fixed physical size increases with decreasing redshift and galaxies undergo intrinsic size growth with decreasing redshift. The implications of fiber aperture bias are investigated in Section~\ref{section:gradients}, where we demonstrate that this has a negligible effect on our conclusions. The conclusion from that section is that we are probing the evolution of the inner $\sim0.3\text{--}3~R_{\rm e}$ of massive quiescent galaxies. We compare the main results shown in Figure~\ref{fig:ages_sdss_mass} with simple conceptual models in an attempt to interpret the trends in the context of the assembly histories of galaxies. We consider four scenarios, graphically represented in Figure~\ref{fig:schematic}. The age and abundance ratio (e.g., [Mg/Fe]) trends with mass are well-known in the local universe. In all four panels, we consider a high-redshift population of quiescent galaxies and trace their evolution over time using their age and abundances, rather than starting with the $z=0$ relation and extrapolating back in time. The stellar population parameters are assumed to be measured within a spectroscopic fiber which has an extent of $\sim1~R_{\rm e}$, unless noted otherwise. The stellar mass spans \mbox{$\approx1$~dex}, the age spans $\approx0.7$~dex, and the abundance ratio spans $\approx0.3$~dex in all panels. At low masses, galaxies are assumed to have solar-scaled abundance ratios. It should be noted that these scenarios are applicable only for $z\lesssim1$. At the highest redshifts, these relations should break down and the [$\alpha$/Fe] abundances are predicted to be super-solar at all stellar masses. This is because a quenched galaxy at very high redshift, regardless of its stellar mass, has to have formed its stars on short timescales to be consistent with the very young age of the universe. The first scenario shows the expected evolution of a population of isolated, passively evolving quiescent galaxies. The galaxies simply age over time, and the trend in abundance ratios remains unchanged. For the next two panels, we assume that a given galaxy doubles its total stellar mass through one or more merger events. The second panel illustrates the evolution of quiescent galaxies undergoing dry (i.e., no star formation) major mergers (1:1) or mass growth of any kind occurring outside the fiber resulting in the doubling of total mass. For comparison, the low-redshift relation for passive evolution is shown as a dotted blue line. The major merger of two similar galaxies results in a more massive galaxy with stellar population properties that are unchanged. This simply moves the new galaxy to the right toward higher masses. This scenario is also equivalent to a passively evolving stellar population in the inner region accompanied by mass buildup in the outskirts. The accumulation of disrupted galaxies and/or in-situ star formation at large radius leads to an increase in total stellar mass and shifts the galaxy to the right, but the spectroscopic data are insensitive to these outer regions and thus stellar population parameters remain unchanged. The third panel shows the evolution of quiescent galaxies undergoing dry minor mergers (1:10) within $\sim1~R_{\rm e}$. The accretion of a lower mass galaxy ``dilutes'' both age and abundances, shifting the new galaxy both downward and to the right from the original relation. For an equal fractional gain in mass (i.e., a factor of two in this case), the changes introduced to both age and abundance ratios are thus more severe for minor mergers compared to major mergers. The fourth scenario is a passive evolution model that includes the addition of recently quenched low-mass galaxies, which are assumed to have solar-scaled abundance ratios. The massive end evolves passively with the universe, but the ages at the low-mass end are lowered due to the presence of recently formed stellar populations. These illustrated relations are only meant to represent a simple schematic picture, since the exact relations and relative slopes depend on the details of the merger history as well as the age and abundance distributions of the new quiescent galaxies. Nevertheless, these simple models highlight the power of the measured abundances combined with the age to probe the evolution of quiescent galaxies. All scenarios produce subtle changes in age, which are difficult to reliably detect owing to the systematic uncertainties affecting stellar age measurements. However, in the abundance-mass space, the changes are more pronounced and qualitatively different, thereby enabling greater distinction between the different scenarios. Our main result is that at fixed stellar mass, the abundance ratios at difference redshifts vary by less than 0.1~dex for most cases. This conclusion is most robust for the massive galaxies ($>10^{10.5}~\msun$) because there are more data spanning a large redshift range at higher masses. The lack of evolution with cosmic time in the mass-abundance ratio correlation leads us to favor scenario I. Taken at face value, these results favor a scenario in which the inner $\sim0.3\text{--}3~R_{\rm e}$ of massive quiescent galaxies have been passively evolving for the past $\sim7$~Gyr. As we demonstrate at the end of this section, the results are also consistent with modest mass growth between $z=0.7$ and $z=0.1$ in the outskirts (scenario II). Due to the absence of lower mass galaxies at high redshifts in our sample, we cannot draw definitive conclusions about the evolution of these objects from the data. However, the young ages of these galaxies combined with their roughly solar-scaled abundances indicate that our results are consistent with the addition of new, low-mass quiescent galaxies with solar-scaled abundance ratios (see scenario IV in Figure~\ref{fig:schematic}). The addition of low-mass quiescent galaxies over time is also favored by simulations \citep[e.g.,][]{Cen2014} as well as independent data, such as evolution in the luminosity and mass functions, which indicate that the number of red galaxies has doubled since $z=1$ \citep[e.g.,][]{Faber2007, Pozzetti2010, Moustakas2013}. As demonstrated in Figure~\ref{fig:logage_redshift}, the derived SSP-equivalent ages are considerably younger than the age of the universe at all epochs, suggesting an \emph{equivalent} single-burst star formation epoch of $z_{\rm f}\lesssim1.5$. We stress that $z_{\rm f}$ is a representative parameter, and not suggestive of the {\it actual} formation epoch. The real SFH is likely more complex and extended in time, but it is a common point of comparison to represent complex SFHs with an SSP-equivalent, or effective single-burst, epoch \citep[e.g.,][]{Treu2005, vanDokkum2007}. In other words, this result should not be interpreted as the galaxies in each mass--redshift bin having uniformly formed their stars in a single burst at $z_{\rm f}\lesssim1.5$. Instead, the low value of $z_{\rm f}$ implies that our results are inconsistent with all of the stars in the galaxies in our sample having formed at very high redshifts \cite[see also the discussion in ][]{Schiavon2006}. The addition of newly quenched galaxies at $z_{\rm f}\gtrsim1$ \citep{vanDokkum2001} naturally explains the young ages of galaxies in our sample, and is supported by both simulations and observations \citep[e.g.,][]{Whitaker2010, ElicheMoral2010, Prieto2013, Cen2014, Marchesini2014}. In order to be consistent with the apparent passive evolution since $z=0.7$, young quiescent galaxies cannot be entering the sample in large numbers at these late times at the highest masses. The inhomogeneous nature of the $z<0.7$ quiescent population suggests that there may be variations in the ages of galaxies within a given mass--redshift bin (see also \citealt{Whitaker2010}). We have tried performing unweighted stacking to ensure that a few bright, young galaxies, which are given more weight during the stacking procedure due to their smaller flux errors, are not driving the ages to low values. The resulting best-fit ages as well as the abundances agree with results from weighted stacking to within 1$\sigma$ errors, thereby demonstrating that this is a negligible effect. In addition, there is no correlation between S/N and $U-V$ color of individual galaxies within each mass--redshift bin, indicating that there is no obvious bias against the reddest galaxies, i.e., the reddest galaxies are equally likely to have large weights in the stacks as the bluest galaxies in the sample. As discussed in Appendix~\ref{section:stacking_test}, the best-fit parameter measured from the stacked spectrum and the weighted average of best-fit parameters from fitting individual spectra agree to within $0.05$~dex. Moreover, the unweighted average of the best-fit ages resulting from fitting individual galaxies agrees with the weighted average age to within 1~Gyr. Thus we conclude that these low ages are representative of the galaxies in the sample, though it should be noted that these are still SSP-equivalent ages. Our results are also consistent with the conclusions from size evolution studies \citep{Daddi2005, Trujillo2006, vanDokkum2008, vanderWel2008, Cimatti2008, Bezanson2009, Damjanov2009, Williams2010, Cassata2010, vanDokkum2010, LopezSanjuan2012, McLure2013, Belli2013}. For example, \cite{vanDokkum2010} found that stellar mass in the central regions (inner 5~kpc) has remained roughly constant with redshift, while the outer regions (out to $\sim75$~kpc) of massive quiescent galaxies have been gradually building up over the last 10~Gyr. In other words, stellar populations have been passively evolving in the centers of massive quiescent galaxies, undisturbed by bursts of star formation or merger activities, while the outskirts have been evolving overtime. Mass and size growth is thought to occur mostly via minor mergers, although in-situ star formation may contribute $\sim20\%$ to the total mass buildup \citep{vanDokkum2010}. As previously discussed, stellar masses quoted in this work nominally represent the total stellar mass. As noted above, the size growth of massive galaxies suggests that they are growing in total mass as well. However, the amount of mass growth from $z=0.7$ to $z=0.1$ is not dramatic, amounting to $\lesssim30\%$ or $\lesssim0.15$ dex \citep{vanDokkum2010}, which is comparable to the uncertainties in our stellar mass estimates. Moreover, since the adopted mass bins for creating stacked spectra are $\gtrsim0.3$ dex wide, the effects of modest mass growth (e.g., 0.15~dex) will not be discernible in the present work. In other words, our main results taken at face value favor a passively evolving stellar population since $z=0.7$, but they are also consistent with modest mass growth in the outskirts over time. \subsection{Comparisons to Previous Work} Our results broadly agree with previous work on stellar populations in quiescent galaxies at low redshift \citep[e.g.,][]{Trager2000b, Saglia2002, Cenarro2003, Eisenstein2003, Thomas2003, Worthey2004, Thomas2005, Sanchez2006b, Graves2007, Schiavon2007, Graves2008, Matkovic2009, Smith2009, Zhu2010, Johansson2012, Worthey2014, Conroy2014}. The primary conclusions from low-redshift studies, including our own, is that more massive galaxies are older, more metal-rich, and have enhanced abundance ratios in the elements Mg, C, and N compared to lower mass galaxies. Owing to the demanding S/N requirements, there have been relatively few studies\footnote{During the refereeing process, we became aware of \cite{Gallazzi2014}, who analyzed optical spectra of $\sim70$ star-forming and quiescent galaxies in the redshift range $0.65\geq z \geq 0.75$ and obtained relations between light-weighted stellar age, stellar mass, and stellar metallicity for both the total galaxy population and for star-forming and quiescent galaxies separately. The authors concluded that both the metallicity and age evolutions of the quiescent galaxies at $z=0.7$ are consistent with passive evolution, in excellent agreement with our results.} of stellar abundances patterns at $z\gtrsim0.1$. \cite{Kelson2006} examined 19 cluster elliptical and lenticular galaxies at $z=0.33$ and measured the age and abundance ratios using eight blue Lick/IDS indices. Interestingly, while they found that total metallicity and [N/Fe] are tightly correlated with velocity dispersion, they did not find significant variations in age, [$\alpha$/Fe], nor [C/Fe] with velocity dispersion. \cite{Sanchez2009} analyzed stacked spectra of 215 red-sequence galaxies in cluster and group environments from $z\sim0.75$ to 0.45 using Lick/IDS indices. The authors confirmed that massive galaxies have $Z\sim \rm Z_{\odot}$ and are enhanced in $\alpha$-elements. Additionally age variation was found to be consistent with passive evolution, all in agreement with the present work. They concluded that massive galaxies formed their stars at high redshift and passively evolved with time, while lower mass galaxies experienced longer star formation episodes and thus have constant luminosity-weighted ages over the redshift range considered. Also in qualitative agreement with our results, they demonstrated that the total metallicity and [$\alpha$/Fe] as a function of velocity dispersion (or mass) are constant with time. \cite{Schiavon2006} went beyond the cluster environment and pushed to even higher redshifts, analyzing Keck DEIMOS spectra of $0.7\leq z \leq 0.9$ red field galaxies. Their stellar population synthesis modeling showed that $z\sim0.9$ galaxies have mean light-weighted ages of only 1~Gyr. The authors concluded that either individual galaxies are experiencing low-level star formation (i.e., ``frosting") or galaxies with younger stars are continually being added to the quiescent population. Interestingly, the time elapsed between $z\sim0.9$ and $z\sim0.7$ is roughly 1 Gyr, and we estimate the youngest $z\sim0.7$ AGES galaxies to be approximately 2--3 Gyr in age. It is thus plausible that the $z\sim0.9$ quiescent galaxies from \cite{Schiavon2006} are progenitors of the $z\sim0.7$ galaxies in our sample. Recently, \cite{Jorgensen2013} analyzed the stellar populations of quiescent galaxies in three galaxy clusters at $z=0.54$, 0.83, and 0.89, and inferred that while the evolution in the fundamental plane is consistent with passive evolution, the variations in total metallicity and [$\alpha$/Fe] with redshift appear to be inconsistent with a passive evolution scenario in the redshift interval considered. Their velocity dispersion-line indices scaling relations indicate that the blue metal lines are stronger than expected for passive evolution. At least some of the discrepancies between conclusions of previous work and our favored interpretation can be attributed to differences in the analysis (e.g., different stellar population synthesis models) and sample selection. The present work selects quiescent galaxies based on sSFR estimated from SED-fitting, but other options for quiescence selection include various combinations of morphological, emission line, S/N, and/or color cuts. In contrast to our quiescent sample which was selected to be mass-complete, most previous high-redshift studies specifically targeted group and cluster environments (of the four discussed above, only \citealt{Schiavon2006} examined field galaxies) due to the efficiency of obtaining many simultaneous spectra with current generation multi-object spectrographs. Typically a few groups or clusters were analyzed at a time, rendering the task of connecting the results to the global quiescent galaxy population difficult as the results can be affected by small-number statistics. These factors make a straightforward and detailed comparison between our results and previous work very challenging. The role of environment on stellar population properties at fixed stellar mass is a subject of ongoing debate in the literature \citep{Sanchez2003, Eisenstein2003, Thomas2005, Sanchez2006b, Trager2008, Toloba2009, Zhu2010, Thomas2010, Johansson2012}. In the present work, we have found that the Keck DEIMOS sample consisting of two bright quiescent galaxies in a $z=0.83$ cluster are older compared to the average quiescent galaxy at $z=0.7$, but their ages are consistent with the age of the universe at $z=0.83$. The results hint at the impact of environment on stellar populations of galaxies, but larger samples will be required to verify these tantalizing trends. \subsection{Radial Gradients and Fiber Size} \label{section:gradients} An important effect that needs to be considered is the fraction of the galaxy observed as a function of both the fiber size and the redshift of the galaxy. The Hectospec instrument on the MMT has a fiber that subtends $1\farcs5$ on the sky, whereas the SDSS spectroscopic fiber has a diameter of $3\farcs$ For reference, $1\farcs5$ corresponds to 2.8~kpc and 10.7~kpc at $z=0.1$ and 0.7, respectively. There are two effects that need to be considered. First, the angular extent of a galaxy of fixed physical size is larger at smaller redshift. Second, quiescent galaxies are believed to have been growing in both size and mass over the last $\sim$10 Gyr \citep[e.g.,][]{Naab2009, vanDokkum2010}. This implies that we are sampling different fractions of the galaxy at different redshifts, both due to the varying angular diameter distance and to the intrinsic growth of the galaxy with redshift. In the presence of radial gradients, this could introduce a bias (with respect to the global properties of the galaxy) in the inferred stellar population parameters. Spatially resolved observations of galaxies at low redshift have enabled radial gradient measurements of various stellar population parameters \citep[e.g.,][]{Gorgas1990, Sanchez2007, Baes2007, Brough2007, Rawle2010, Spolaor2010, Loubser2012, LaBarbera2012, Greene2013}. \cite{Greene2013} measured the radial gradients of a sample of 33 massive quiescent galaxies and found a modest negative gradient in [Fe/H] and a strong radial decline in [C/Fe], but almost constant [Mg/Fe], [N/Fe], and [Ca/Fe] out to 2.5~$R_{\rm e}$. Additionally, the age gradient was measured to be very weak or nonexistent. Similarly, \cite{Spolaor2010} inferred almost no radial gradient in age and [$\alpha$/Fe] but a weak negative gradient in [Z/H] from a sample of 37 massive quiescent galaxies. From these local studies it thus appears unlikely that our age and [$\alpha$/Fe] measurements are biased by the sampling radius of the fiber, but we cannot immediately rule out the possibility that our [Fe/H] and [C/Fe] measurements are affected by radial gradients. \begin{figure} \centering \subfigure{ \includegraphics[width=0.9\columnwidth]{RfibRe_vs_z_diffmasses_ages_sdss-crop.pdf} } \subfigure{ \includegraphics[width=0.9\columnwidth]{apparent_evol_onegradient-crop.pdf} } \caption{{\it Top panel:} the ratio of fiber radius to the effective radius as a function of redshift. The effective radius is computed by assuming a mass-size relation at $z\sim0$ from SDSS data \citep{Shen2003} and applying a redshift evolution scaling $R_{\rm e} \propto (1+z)^{-1}$ \citep[e.g.,][]{Williams2010}. Each curve spans the entire redshift range sampled by the data in each mass bin. The dashed and solid lines correspond to SDSS and AGES data, respectively. The mass bins increase from top to bottom (gray to pink). The mid-bin values for each mass interval are $\log (M/\msun) = 9.80$, 10.15, 10.45, 10.75, 11.05, and 11.35 for AGES, and $\log (M/\msun) = 10.58$, 10.93, 11.28, and 11.63 for SDSS. {\it Bottom panel:} the differences in inferred [Fe/H] between $z=0.1$ and five different values of $z$ ranging from 0.2 to 0.7, shown in different colored curves for a metallicity gradient of $\nabla_{\rm [Fe/H]} \equiv{\rm \Delta [Fe/H]}/\Delta \log {(R/R_{\rm e})}=-0.23$~dex~decade$^{-1}$. The black circles represent the maximum aperture bias between the lowest and highest redshift bins within each AGES stellar mass bin.} \label{fig:gradients} \end{figure} We carry out simple simulations to quantitatively assess this possibility. We investigate [Fe/H] only, but similar conclusions can be drawn for [C/Fe]. We adopt a metallicity gradient from \cite{Spolaor2010}, corresponding roughly to $\nabla_{\rm [Fe/H]} \equiv{\rm \Delta [Fe/H]}/\Delta \log {(R/R_{\rm e})}=-0.23$~dex~decade$^{-1}$ over the range $0.5~R_{\rm e}$--$2.5~R_{\rm e}$, assuming that [Z/H]~$\propto$~[Fe/H]. The metallicity profile ${\rm [Fe/H]}(r)$ is convolved with an $n=4$ Sersic profile $I(r)$ to obtain a light-weighted average metallicity within a given radius, $\langle {\rm [Fe/H]}(r) \rangle$. The effects of seeing ($\approx1\farcs5$ for AGES, which is comparable to the diameter of the fiber) is ignored because while light is scattered out of the fiber, light from larger radius is also scattered inward. Thus light captured by the fiber is actually sampled from a larger radius, on average, thereby effectively reducing the magnitude of aperture bias. Here we define $f_{\rm origin}$, the fraction of light originating from outside the fiber radius that is scattered into the fiber, i.e., $f_{\rm origin}=0$ in the absence of seeing. In the high-$z$ limit where the fiber size and seeing are larger than the effective radius, suppose by a fiducial factor of three, $f_{\rm origin}~\approx0.1$. In the low-$z$ limit where the effective radius is roughly three times larger than both the fiber and the seeing, $f_{\rm origin} \approx0.3$, which is non-negligible. In the calculations below we do not include the effects of seeing because it provides the worst-case scenario for bias effects. The bias introduced as a result of observing within an aperture of radius $r$ is calculated following \cite{Moustakas2011}: \begin{equation} \Delta \langle {\rm [Fe/H]}(r) \rangle = \langle {\rm [Fe/H]}(r) \rangle - \langle {\rm [Fe/H]}_{\infty} \rangle \;, \end{equation} where \begin{equation} \langle {\rm [Fe/H]}(r) \rangle = \frac{\int_{0}^{r}{\rm [Fe/H]}(r^\prime)I(r^\prime)r^\prime dr^\prime}{\int_{0}^{r} I(r^\prime)r^\prime dr^\prime}\;, \end{equation} and $\langle {\rm [Fe/H]}_{\infty} \rangle$ is the same quantity integrated out to $\infty$. For $\nabla_{\rm [Fe/H]} = -0.23$~dex~decade$^{-1}$, the bias introduced is relatively modest, corresponding to $\approx 0.08$~dex and $\approx 0.05$~dex when probing $0.5~R_{\rm e}$ and $2~R_{\rm e}$, respectively. When the gradient is twice as strong, akin to that found by \cite{Greene2013}, the resulting bias is roughly doubled. The goal of this section is to investigate how much apparent evolution in [Fe/H] in our sample could be due to the effects of aperture bias. We estimate the average sizes of our quiescent galaxies as a function of redshift (direct size estimates are unfortunately not available for individual galaxies) to determine the fraction of the galaxy probed by the spectroscopic fiber as a function of redshift. To do this we assume a mass-size relation at $z\sim0$ from SDSS data \citep{Shen2003} and apply a redshift evolution scaling $R_{\rm e} \propto (1+z)^{-1}$ \citep[e.g.,][]{Williams2010}. The top panel of Figure~\ref{fig:gradients} shows the ratio of fiber radius to the effective radius as a function of redshift. Next, we compute the differences in inferred [Fe/H] between $z=0.1$ and five different values of $z$ ranging from 0.2 to 0.7 purely due to the sampled fraction of the galaxy evolving with redshift, shown as different colored curves in the bottom panel of Figure~\ref{fig:gradients}. The black circles represent the maximum aperture bias between the lowest and highest redshift bins within each AGES stellar mass bin. At lower masses, an apparent evolution of $\approx0.06$~dex is induced due to the evolving aperture bias, implying that for an intrinsically unevolving population, we would observe [Fe/H] to increase by $\approx0.06$~dex with time in the redshift range probed (e.g., from $z=0.4$ to $z=0.1$ for the third lowest mass bin). Apparent evolution is the strongest for low-mass galaxies, and the effect becomes weaker with increasing mass. If radial gradients in [Fe/H] are present in our sample at the level we have assumed for this test, then this apparent evolution can masquerade as a true intrinsic evolution in stellar population parameters purely as a consequence of probing different fractions of the galaxy. However this is a small effect, amounting to at most $\lesssim0.06$~dex in our sample (or $\lesssim0.1$~dex in the more extreme case where the gradient is twice as strong for our sample of quiescent galaxies). Moreover, if we account for the effects of seeing, the magnitude of the bias would be even smaller. In comparison, the observed variation in the stellar population parameters as a function of redshift is $\lesssim0.1$~dex (see Figure~\ref{fig:ages_sdss_mass}). Thus we conclude that aperture bias has a minor effect on the interpretation of our results. In this paper we measured SSP-equivalent ages and detailed abundance patterns of quiescent galaxies using stacked SDSS and AGES spectra and individual spectra of two brightest $z=0.83$ cluster galaxies from \cite{Holden2010}. The main sample spans a redshift interval of $0.1<z<0.7$ and a stellar mass range from $10^{9.6}$ to $10^{11.8}~\msun$. We selected quiescent galaxies based on their star formation rates estimated from SED-fitting. The AGES sample of quiescent galaxies were divided into five redshift intervals each spanning roughly 1~Gyr in cosmic time and further divided into mass bins. The mass bins were chosen such that the sample was complete in stellar mass at each redshift. The stacked spectra were fit using a full spectrum modeling MCMC code developed by \cite{Conroy2012}. We also carried out a variety of systematic tests to examine the robustness of the modeling code. We now summarize our main results. \begin{enumerate} \item We confirm earlier results of stellar population modeling of quiescent galaxies at low redshift. Massive galaxies harbor old stellar populations with roughly solar [Fe/H] abundances and enhanced [Mg/Fe], [C/Fe], and [N/Fe] abundance ratios. Adopting [Mg/Fe] as a star-formation timescale indicator, massive galaxies seem to form their stars on shorter timescales compared to lower mass galaxies. \item There is negligible evolution in the abundances of Fe, Mg, C, N, and Ca at fixed stellar mass over roughly 7~Gyr of cosmic time, and the evolution of the stellar ages of massive galaxies is consistent with passive evolution since $z=0.7$. The 0.1~dex or smaller variation in abundance ratios (e.g., [Fe/H], [Mg/Fe], [Ca/Fe]) as a function of stellar mass from $z=0.1$ to $z=0.7$ puts a stringent constraint on the assembly histories of galaxies over the redshift interval considered. Our results support the passive evolution of the inner $\sim0.3\text{--}3\;R_{\rm e}$ of massive quiescent galaxies ($M>10^{10.5}~\msun$) over the last $\sim7$ Gyr. At lower masses our results are also consistent with the addition of younger, newly quenched galaxies over time. \item The derived SSP-equivalent ages are considerably younger than the age of the universe at all epochs, consistent with an \emph{equivalent} single-burst star formation epoch of $z_{\rm f}\lesssim1.5$. The addition of newly quenched galaxies at $z_{\rm f}\gtrsim1.5$ naturally explains the young ages of galaxies in our sample. These young stellar population ages coupled with the existence of massive quiescent galaxies at $z>1$ indicate the inhomogeneous nature of the $z\lesssim0.7$ quiescent population. In order to be consistent with the apparent passive evolution since $z=0.7$, young quiescent galaxies cannot be entering the sample at these late times at the highest masses. \item There is tentative evidence that galaxies in cluster environments are older than galaxies inhabiting low-density environments. We hesitate to draw any strong conclusions, however, due to the small sample size as well as the possibility of contamination from atmospheric absorption and sky emission features imparting systematic effects on our spectral modeling results. This stresses the importance of obtaining deeper spectra with careful attention to telluric corrections and sky subtractions in addition to the development of tools capable of modeling low S/N spectra of these high-redshift galaxies. \item The full spectrum fitting approach allows reliable and accurate abundance measurements, including age, Fe, Mg, C, N, and Ca, down to low S/N. The ability to accurately measure detailed abundance patterns in low S/N spectra with reliable uncertainty estimates opens the possibility to engage in detailed stellar population analysis at high redshift and in the low surface brightness outskirts of nearby galaxies. \end{enumerate}	14	3	1403.4932	We present results from modeling the optical spectra of a large sample of quiescent galaxies between 0.1 < z < 0.7 from the Sloan Digital Sky Survey (SDSS) and the AGN and Galaxy Evolution Survey (AGES). We examine how the stellar ages and abundance patterns of galaxies evolve over time as a function of stellar mass from 10<SUP>9.6</SUP>-10<SUP>11.8</SUP> M <SUB>⊙</SUB>. Galaxy spectra are stacked in bins of mass and redshift and modeled over a wavelength range from 4000 Å to 5500 Å. Full spectrum stellar population synthesis modeling provides estimates of the age and the abundances of the elements Fe, Mg, C, N, and Ca. We find negligible evolution in elemental abundances at fixed stellar mass over roughly 7 Gyr of cosmic time. In addition, the increase in stellar ages with time for massive galaxies is consistent with passive evolution since z = 0.7. Taken together, these results favor a scenario in which the inner ~0.3-3 R <SUB>e</SUB> of massive quiescent galaxies have been passively evolving over the last half of cosmic time. Interestingly, the derived stellar ages are considerably younger than the age of the universe at all epochs, consistent with an equivalent single-burst star formation epoch of z <~ 1.5. These young stellar population ages coupled with the existence of massive quiescent galaxies at z > 1 indicate the inhomogeneous nature of the z <~ 0.7 quiescent population. The data also permit the addition of newly quenched galaxies at masses below ~10<SUP>10.5</SUP> M <SUB>⊙</SUB> at z < 0.7. Additionally, we analyze very deep Keck DEIMOS spectra of the two brightest quiescent galaxies in a cluster at z = 0.83. There is tentative evidence that these galaxies are older than their counterparts in low-density environments. In the Appendix, we demonstrate that our full spectrum modeling technique allows for accurate and reliable modeling of galaxy spectra to low S/N (~20 Å<SUP>-1</SUP>) and/or low spectral resolution (R ~ 500).	false	[ "massive quiescent galaxies", "quiescent galaxies", "massive galaxies", "galaxies", "stellar population synthesis", "cosmic time", "stellar ages", "z", "stellar mass", "fixed stellar mass", "time", "Evolution Survey", "lt", "low spectral resolution", "R", "lt;~", "These young stellar population ages", "mass", "abundance patterns", "elemental abundances" ]	12.285433	6.699439	189
514111	[ "Smiljanic, R." ]	2014Ap&SS.354...55S	[ "Stellar abundances of beryllium and CUBES" ]	6	[ "Department for Astrophysics, Nicolaus Copernicus Astronomical Center, Toruń, Poland" ]	[ "2020past.conf...29S", "2021A&A...646A..70S", "2021FrASS...8....6R", "2023ExA....55....1E", "2023ExA....55...95S", "2023ExA....55..117G" ]	[ "astronomy" ]	1	[ "Globular clusters: general", "Stars: abundances", "Stars: late-type", "Stars: population II", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1895ApJ.....1...14R", "1929ApJ....70...11R", "1954ApJ...119..113G", "1955ApJS....2..167F", "1955PhRv...97.1237S", "1967AnPhy..44..426B", "1967ApJ...148....3W", "1968ApL.....2..235H", "1968PASP...80..622W", "1968SoPh....5..159G", "1969ARA&A...7...99W", "1970Natur.226..727R", "1971A&A....15..337M", "1971ApJ...164..111C", "1973ApJ...185L..27B", "1974SoPh...36...11R", "1975A&A....40...99M", "1975A&A....42...37C", "1976PASP...88..353B", "1977ApJ...214..124B", "1979AJ.....84.1756S", "1982A&A...115..357S", "1984A&A...139..394M", "1988A&A...193..193R", "1990ApJ...364..568V", "1992ApJ...388..184R", "1992ApJ...401..584D", "1992Natur.357..379G", "1993AJ....106.2309B", "1993ApJ...403..630P", "1993ApJ...406..569T", "1995A&A...302..184G", "1995ApJ...447..184Y", "1996A&A...313..909G", "1997A&A...319..593M", "1997A&A...321L..37D", "1997ApJ...480..784P", "1997ApJ...488..515O", "1997ApJ...488..730R", "1997JPCRD..26.1185K", "1998A&A...337..714V", "1998ApJ...499..735L", "1998Natur.392..791B", "1999A&A...345..249C", "1999AJ....117.1549B", "1999ApJ...511..466B", "2000A&A...362..666P", "2000A&A...364L..42P", "2000IAUS..198..425B", "2000SPIE.4008..534D", "2001A&A...369...87G", "2001ApJ...546L..65B", "2001ApJ...549..303S", "2002ApJ...566..252V", "2004A&A...417..769A", "2004A&A...426..651P", "2005A&A...436L..57P", "2005A&A...439..227M", "2005ARA&A..43..481A", "2005PhyS...72..309K", "2006A&A...447..299G", "2006ApJ...641.1122B", "2006GGRes..30..245M", "2007A&A...464..601P", "2007A&A...464.1029D", "2008ApJ...673..676R", "2008ApJ...681...18K", "2009A&A...499..103S", "2009A&A...499..835V", "2009A&A...505..139C", "2009ARA&A..47..481A", "2009ApJ...698L..37I", "2009ApJ...701.1519R", "2009MNRAS.392..205T", "2010A&A...510A..50S", "2010A&A...524L...2S", "2010ASSP...16..379L", "2011A&A...535A..75S", "2011A&A...536A.105V", "2011ApJ...738L..33T", "2011ApJ...743..140B", "2012A&A...542A..67P", "2012A&ARv..20...50G", "2012ASPC..458...79S", "2012ApJ...744..158C", "2013ApJ...773...33I", "2013MNRAS.428..236O", "2014A&A...563A...3P", "2014ApJ...784..153H" ]	[ "10.1007/s10509-014-1916-9", "10.48550/arXiv.1403.6276" ]	1403	1403.6276_arXiv.txt	\label{sec:intro} Beryllium has one single stable isotope, $^{9}$Be. Spectral lines of Be have been identified in the Sun, in the near UV region that is observable from the ground, quite some time ago by \citet{1895ApJ.....1...14R}. The lines tentatively identified included the Be I lines at $\lambda$ 3321.011, 3321.079 and 3321.340 \AA, and the Be II resonance lines at $\lambda$ 3130.422 and 3131.067 \AA\ \citep[wavelengths from][]{1997JPCRD..26.1185K,2005PhyS...72..309K}. However, it has been shown that the identification of the Be I lines was erroneous \citep{1968SoPh....5..159G,1968ApL.....2..235H}. Only the Be II $^{2}$S--$^{2}$P$_{0}$ resonance lines at $\sim$ 3130 \AA\ are now used to determine Be abundances in late-type stars. Other Be lines can be found in the UV below 3000 \AA, but are only observable from space and seem to be of limited use \citep{1996A&A...313..909G}. Isotopic shifted lines of unstable $^{7}$Be and $^{10}$Be have been searched for, but were never detected \citep{1968PASP...80..622W,1973ApJ...185L..27B,1975A&A....42...37C}. High signal-to-noise (S/N) spectra in the near UV is hard to obtain because of atmospheric extinction. In addition, the spectral region of the Be lines is crowded with other atomic and molecular lines (see the solar spectrum in Fig. \ref{fig:sun}). Some of these blending lines are still unidentified, as for example the one contaminating the blue wing of the 3131.067 \AA\ line \citep[see e.g.][]{1974SoPh...36...11R,1995A&A...302..184G,1997ApJ...480..784P,2011A&A...535A..75S}. Early analyses of Be abundances were reviewed by \citet{1969ARA&A...7...99W} and \citet{1976PASP...88..353B}. The solar Be abundance seems to have been first determined by \citet{1929ApJ....70...11R}. The approximation of local thermodynamic equilibrium (LTE) has been shown to result in correct abundances for the Sun \citep{1975A&A....42...37C,1979AJ.....84.1756S,2005ARA&A..43..481A}. The solar Be abundance also seems to be largely insensitive to 3D hydrodynamical effects \citep{2004A&A...417..769A,2009ARA&A..47..481A}. The current value of the solar meteoritic abundance of Be is A(Be)\footnote{The abundance of an element X in this notation is A(X) = $\log \epsilon$(X) = $\log$ [N(X)/N(H)] + 12, i.e. an abundance by number in a scale where the number of hydrogen atoms is 10$^{12}$.} = 1.32 \citep{MakishimaNakamura06,2010ppc..conf..379L} while the photospheric abundance is A(Be) = 1.38 \citep{2009ARA&A..47..481A}. Nevertheless, most 1D LTE model atmosphere analyses of Be in the Sun tend to find values between A(Be) = 1.10 and 1.15. The difference is ascribed to near-UV continuum opacity missing in the computations \citep{1998Natur.392..791B,2001ApJ...546L..65B}. \begin{figure}[t] \begin{center} \includegraphics[width=7cm]{sun} \end{center} \caption{The Be lines in the solar spectrum} % \label{fig:sun} \end{figure} Beryllium is not produced in significant amounts by the primordial nucleosynthesis, because there are no stable elements with mass number 5 or 8 to act as an intermediate step in synthesising $^{9}$Be \citep{1967ApJ...148....3W,1993ApJ...406..569T,2012ApJ...744..158C}. In addition, Be is rapidly destroyed by proton capture reactions when in regions inside a star with a temperature above $\sim$ 3.5 $\times 10^{6}$ K \citep{1954ApJ...119..113G,1955PhRv...97.1237S}. Therefore, stars do not produce Be. In stars like the Sun, for example, Be is only present in the external regions of lower temperature \citep[see Fig.\ 1 of][]{1999ApJ...511..466B}. In evolved stars, where the convective envelope has increased in size and mixed the surface material with the interior, Be, unlike Li, is usually not detected \citep{1977ApJ...214..124B,2006A&A...447..299G,2010A&A...510A..50S}. Beryllium has never been detected in Li-rich giants \citep{1997A&A...321L..37D,1999A&A...345..249C,2005A&A...439..227M}. Long ago, it was understood that Be can only be produced by the spallation of heavier nuclei, mostly from carbon, nitrogen, and oxygen \citep{1955ApJS....2..167F,1967AnPhy..44..426B}. The only known way to produce significant amounts of Be is by cosmic-ray induced spallation in the interstellar medium (ISM), as first shown by \citet{1970Natur.226..727R} and \citet{1971A&A....15..337M}. Two channels of cosmic-ray spallation might work to produce Be. In the so-called \emph{direct process}, Be is produced by accelerated protons and $\alpha$-particles that collide with CNO nuclei of the ISM \citep{1975A&A....40...99M,1990ApJ...364..568V,1993ApJ...403..630P}. In the \emph{inverse process}, accelerated CNO nuclei collide with protons and $\alpha$-particles of the ISM \citep{1997ApJ...488..730R,1998ApJ...499..735L,1998A&A...337..714V,2002ApJ...566..252V,2012A&A...542A..67P}. If the first channel dominates in the early Galaxy, Be should behave as a secondary element, as its production rate would be proportional to the metallicity of the ISM. If the second channel dominates, Be should instead behave as a primary element. The two behaviors can be distinguished by the analysis of Be abundances in metal-poor stars. For primary elements, one should observe a linear correlation between $\log$(Be/H) and the metallicity [Fe/H]\footnote{[A/B] = log [N(A)/N(B)]$_{\rm \star}$ $-$ log [N(A)/N(B)]$_{\rm\odot}$} with a slope close to one. For secondary elements, the slope should be around two. \begin{figure}[t] \begin{center} \includegraphics[width=7cm]{full_box.jpg} \end{center} \caption{The Be abundances as a function of metallicty in the metal-poor stars analyzed by \citet{2009A&A...499..103S}. Also shown are the slopes expected for primary and secondary elements. Clearly, the inverse process dominates the Be production and it behaves as a primary element} % \label{fig:slopes} \end{figure} Linear relations with the two slopes are compared in Fig. \ref{fig:slopes}. Also shown are the Be abundances of metal-poor stars determined by \citet{2009A&A...499..103S}. It is clear then that Be behaves as a primary element, meaning that the inverse process dominates \citep[as first suggested by][]{1992ApJ...401..584D}. The determination of abundances of Be in metal-poor stars were first attempted by \citet{1984A&A...139..394M} and \citet{1988A&A...193..193R}. But it was with the works of \citet{1992ApJ...388..184R} and \citet{1992Natur.357..379G} that the linear correlation with slope of one became well established. Beryllium abundances in metal-poor stars have been further investigated in several works since then \citep[e.g.][]{1993AJ....106.2309B,1997A&A...319..593M, 1999AJ....117.1549B,2009A&A...499..103S,2009MNRAS.392..205T,2009ApJ...701.1519R,2011ApJ...738L..33T,2011ApJ...743..140B}. With the inverse process dominating, and considering that cosmic-rays are globally transported across the Galaxy, then the Be production in the early Galaxy should be a widespread process. It follows that, at a given time, the abundance of Be across the Galaxy should have a smaller scatter than the products of stellar nucleosynthesis (such as Fe and O), as suggested by \citet{2000IAUS..198..425B} and \citet{2001ApJ...549..303S}. In this case, Be abundances could be used as a time scale for the early Galaxy \citep{2005A&A...436L..57P,2009A&A...499..103S}. This is an interesting application that still needs to be further constrained.	Abundances of the fragile element Be can be used in different areas of astrophysics, including studies of the Galactic chemical evolution, of stellar evolution, and of the formation of globular clusters. Only the Be II resonance lines at $\lambda$ 3130.422 and 3131.067 \AA\ are used for abundance studies. The lines are in the near UV spectral region, a region strongly affected by atmospheric absorption. It is very time demanding to obtain the high-resolution, high-S/N spectra needed to study Be. Because of that, Be abundances have not been used to its full potential yet. To increase the number of stars where Be abundances can be determined, and to decrease the time cost of obtaining high-quality spectra, new UV sensitive instruments are needed. CUBES is one such instrument. It has an expected sensitivity gain of about three magnitudes at 3200 \AA\ over UVES. In the near future, CUBES will likely be the only instrument offering the opportunity to measure Be abundances in samples of extremely metal-poor stars and in turn-off stars of globular clusters. Here, preliminary simulations of CUBES-like spectra for these types of stars were presented. For the first case, the simulations indicate that with CUBES spectra it will be possible to investigate the suggested flattening of the relation between Be and Fe at low metallicities ([Fe/H] $<$ $-$3.00). This flattening might have implications, for example, to primordial nucleosynthesis models. With CUBES, it will also be possible to investigate Be abundances in turn-off stars of globular clusters. Beryllium can be used as a tracer of the fraction of polluting material accreted by the second generation stars. If stars with different levels of pollution are found to have the same Be abundances, this would argue against the currently favored dilution-pollution scenarios invoked to explain the chemical properties of globular clusters.	14	3	1403.6276	Stellar abundances of beryllium are useful in different areas of astrophysics, including studies of the Galactic chemical evolution, of stellar evolution, and of the formation of globular clusters. Determining Be abundances in stars is, however, a challenging endeavor. The two Be II resonance lines useful for abundance analyses are in the near UV, a region strongly affected by atmospheric extinction. CUBES is a new spectrograph planned for the VLT that will be more sensitive than current instruments in the near UV spectral region. It will allow the observation of fainter stars, expanding the number of targets where Be abundances can be determined. Here, a brief review of stellar abundances of Be is presented together with a discussion of science cases for CUBES. In particular, preliminary simulations of CUBES spectra are presented, highlighting its possible impact in investigations of Be abundances of extremely metal-poor stars and of stars in globular clusters.	false	[ "atmospheric extinction", "globular clusters", "Stellar abundances", "stellar abundances", "abundance analyses", "abundances", "stellar evolution", "fainter stars", "stars", "CUBES spectra", "CUBES", "current instruments", "the near UV spectral region", "science cases", "Stellar", "different areas", "investigations", "the Galactic chemical evolution", "Galactic", "studies" ]	7.551514	9.321314	-1
746025	[ "Evrard, August E.", "Arnault, Pablo", "Huterer, Dragan", "Farahi, Arya" ]	2014MNRAS.441.3562E	[ "A model for multiproperty galaxy cluster statistics" ]	65	[ "Department of Physics and Michigan Center for Theoretical Physics, University of Michigan, Ann Arbor, MI 48109, USA; Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA; Institut d'Astrophysique, 98bis Bd. Arago, F-75012 Paris, France", "Department of Physics and Michigan Center for Theoretical Physics, University of Michigan, Ann Arbor, MI 48109, USA; École Normale Supérieure de Cachan, 61 Avenue du Président Wilson, F-94230 Cachan, France", "Department of Physics and Michigan Center for Theoretical Physics, University of Michigan, Ann Arbor, MI 48109, USA", "Department of Physics and Michigan Center for Theoretical Physics, University of Michigan, Ann Arbor, MI 48109, USA" ]	[ "2015MNRAS.446.2629E", "2015MNRAS.450.1150E", "2015MNRAS.450.3675S", "2015MNRAS.452.1982W", "2015MNRAS.454.2305S", "2015NatCo...6.7211.", "2016ApJ...831..135N", "2016ApJS..224....1R", "2016MNRAS.455.2149S", "2016MNRAS.456.4020M", "2016MNRAS.460.1270D", "2016MNRAS.460.3900F", "2016MNRAS.461..248S", "2016MNRAS.463..820Z", "2017A&A...604A..89P", "2017ApJ...835..106N", "2017ApJ...844..101A", "2017MNRAS.466.3103S", "2017MNRAS.468.3347S", "2017MNRAS.469.4899M", "2018A&A...614A..72V", "2018MNRAS.474.4089T", "2018MNRAS.474.5500O", "2018MNRAS.477.4931H", "2018MNRAS.478.2618F", "2018RPPh...81a6901H", "2019A&A...632A..54S", "2019MNRAS.482.1352M", "2019MNRAS.484...60M", "2019MNRAS.484.1946G", "2019MNRAS.485.4863M", "2019MNRAS.490.2380C", "2019MNRAS.490.3341F", "2019NatCo..10.2504F", "2020MNRAS.493.1361F", "2020MNRAS.493.4591P", "2020MNRAS.495..428C", "2020MNRAS.495..686A", "2020MNRAS.498.5450P", "2020PhRvD.102b3509A", "2021A&A...653A.135O", "2021AJ....161...30F", "2021MNRAS.501.2467S", "2022A&A...661A..11C", "2022ApJ...926...45N", "2022ApJ...931..166F", "2022ApJ...933...48F", "2022MNRAS.509.3441A", "2022MNRAS.510.2980A", "2022MNRAS.515.4471W", "2022MNRAS.515.4722H", "2022PhR...973....1D", "2022hxga.book...65L", "2023MNRAS.524.3289H", "2023arXiv231012528H", "2023arXiv231211789Z", "2024A&A...685A..61B", "2024ApJ...966..227C", "2024MNRAS.527.9378A", "2024MNRAS.530.3127Z", "2024MNRAS.531.1685N", "2024MNRAS.tmp.1471P", "2024OJAp....7E..29S", "2024arXiv240112049E", "2024arXiv240603180K" ]	[ "astronomy" ]	7	[ "galaxies: clusters: general", "cosmology: theory", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "2001MNRAS.321..372J", "2001MNRAS.323....1S", "2002ApJ...573....7E", "2005A&A...441..893A", "2005RvMP...77..207V", "2007ApJ...660..239K", "2007ApJ...668..772M", "2008ApJ...672..122E", "2008ApJ...675.1106R", "2008ApJ...688..709T", "2009A&A...498..361P", "2009ApJ...691.1307H", "2009ApJ...692.1033V", "2009MNRAS.394L..11S", "2010ApJ...708..645R", "2010ApJ...715.1508S", "2010MNRAS.404.1922A", "2010MNRAS.406.1759M", "2010MNRAS.408.1818W", "2011A&A...526A.105Z", "2011A&A...536A..12P", "2011ARA&A..49..409A", "2011ApJ...732..122B", "2011MNRAS.416..801F", "2011PhRvD..83b3015M", "2012A&A...539A.120B", "2012ApJ...746..178R", "2012MNRAS.425.2244B", "2012MNRAS.426.1829N", "2012MNRAS.426.2046A", "2012NJPh...14e5018R", "2013A&C.....3...23M", "2013ApJ...763..147B", "2013ApJ...777..123B", "2013JCAP...10..060S", "2013MNRAS.430.2638M", "2013MNRAS.435.1265E", "2013PhR...530...87W", "2013SSRv..177..247G", "2014A&A...564A.129I", "2014ApJ...782..107N", "2014ApJ...783...80R", "2014ApJ...785..104R", "2014MNRAS.437.1171M", "2014MNRAS.438...49R", "2014MNRAS.438...62R", "2014MNRAS.438...78R", "2014MNRAS.439..588B", "2014MNRAS.439.2485C", "2014MNRAS.440.2115J", "2014MNRAS.440.2290M", "2014MNRAS.441.1270L" ]	[ "10.1093/mnras/stu784", "10.48550/arXiv.1403.1456" ]	1403	1403.1456_arXiv.txt	Counts of galaxy clusters provide constraints on cosmological parameters \citep[\eg][and references therein]{Voit0410173, Allen1103.4829}, and test fundamental theories of gravity and cosmic acceleration \citep[\eg][]{Weinberg1201.2434}. Such studies typically use cluster samples identified via optical \citep{Rozo10}, X-ray \citep{Mantz0909.3098, Henry0809.3832} or thermal Sunyaev-Zel'dovich \citep[SZ,][]{Benson1112.5435, Sievers1301.0824, PlanckXX1303.5080} signatures of the baryons in the halos that host cluster phenomena. These analyses are empowered by simulation studies that calibrate the space density as a function of halo mass, known as the \emph{mass function}, within a given cosmology \citep{Tinker08, Bhattacharya1005.2239, Murray1306.6721}. Modeling the expected counts of clusters in a wide-area observational survey requires combining the mass function with a statistical model that expresses the likelihood for a halo of mass $M$ at redshift $z$ to have some intrinsic observable signal, $S$, detectable in the survey. Evidence from observations \citep{Arnaud0502210, Maughan0703504, Pratt0809.3784, Vikhlinin0805.2207, Zhang1011.3018, Ruel1311.4953, Saliwanchik1312.3015, Ettori1307.7157, Maughan1212.0858} and simulations \citep{Evrard08, Stanek0910.1599, Fabjan1102.2903, Munari1301.1682, Jiang1311.6649, LeBrun1312.5462, Biffi1401.2992} support a model in which the scaling law behavior is power-law with mass in the mean, with approximately log-normal variance. While scaling behavior of cluster properties has been studied for decades \citep[see][for a recent review]{Giodini1305.3286}, most works have focused on correlating pairs of observed signals, $\{ S_2, S_1 \}$, or on studying how a single observable scales with mass, $\{S_1, M \}$. Simulations provide a natural environment for the latter, since the true halo mass is known. For observations, mass estimates are made indirectly from measured signals, for example through assumption of virial or hydrostatic equilibrium, and this methodology introduces the potential for bias and additional variance that must be calibrated \citep[\eg][]{Rasia1201.1569, Battaglia1209.4082, Nelson1308.6589}. Alternatively, masses can be inferred through inversion of a given observable-mass relation. In this way, an observable serves as proxy for halo mass. Evidence of biases in mass proxies can arise from comparisons among different observable signals. Planck satellite measurements of the thermal SZ effect in the optically-selected \maxBCG\ sample \citep{Planck11_optical} led to a detailed re-examination of X-ray, SZ and optical scaling relations by \citet{RozoI,RozoII,RozoIII}. That study concluded that the Planck $\Ysz$ mass calibration was biased low by a few tens of percent, a finding supported by independent weak gravitational lensing estimates of Planck clusters \citep{VonderLinden1402.2670}, although other studies are less supportive \citep{Israel1402.3267}. \citet{RozoII} present a model for multivariate signal counts and other statistics under the assumption of a locally power-law mass function. That model was employed to interpret a combined set of X-ray, SZ and optical data, resulting in a set of preferred scaling relations presented in \citet{RozoIII} (see their Table~4). In this paper, we present a non-local extension of that model that expands its scope to effectively cover the complete dynamic range of properties displayed by the population of galaxy clusters. Within a $\Lambda$CDM cosmology, we show that the mass function of the massive halos that host groups and clusters of galaxies can be represented by a low-order polynomial (in log-space) to a typical accuracy of a few percent, comparable to its calibrated level of precision from N-body simulations. Convolving this mass function representation with a multivariate Gaussian of logarithmic halo properties at fixed mass and redshift results in analytic expressions for the space density as a function of multiple observables and other derivative statistics. By offering a fast method for estimating survey sample and follow-up study outcomes, this formalism is intended to complement data analysis methods based on similar model assumptions \citep{Maughan1212.0858}. We employ a halo mass convention of $\Mfiveh$, the mass within a spherical region encompassing $500$ times the critical density, $\rho_c(z)$, but the analytic expressions can be applied using any choice of halo mass convention. We use $\Mfiveh$ to be consistent with the scaling laws presented in \citet{RozoIII}. In \S\ref{sec:model}, we derive expressions for multi-observable cluster population statistics using low-order polynomial approximations of the mass function in log space. These are then applied to X-ray and SZ statistics in \S\ref{sec:applications}. In \S\ref{sec:discussion}, we discuss some of the model's strengths and limitations, and we summarize our results in \S\ref{sec:summary}	\label{sec:discussion} We have emphasized an application of the model to the massive halos that host galaxy clusters at late cosmic times, but the mathematical framework is general, so the model could be applied at earlier epochs to describe phenomenology associated with the high-mass end of the mass function. Galaxies and quasars at redshifts of a few or early star formation at redshifts of tens are potential applications. The key requirements are observables or properties that of halos that scale as power-laws with halo mass in the mean, with variability described by a log-normal covariance. Applied to groups and clusters, the third-order model is essentially global in scope. Compared to local Tinker convolution estimates, the cubic approximation achieves better than $10\%$ accuracy over nearly the whole signal ranges covered by current observations, and for redshifts $z < 1.5$. This level of accuracy is comparable to the current level of systematic uncertainty in the mass function derived from simulations, particularly when the effects of baryon physics are included \citep{Stanek0809.2805, Martizzi1307.6002, Cusworth1309.4094, Cui1402.1493}. The pivot location sets the range of accuracy for the lower-order approximations. Since the third-order approximation is based on a Taylor expansion of $e^{-\betathree \mu^3/6} \simeq 1 - \betathree \mu^3/6$, where $\mu = \ln(M/M_p)$, the model breaks down as $\mu^3 \rightarrow 6/\betathree$. For the pivots chosen here, this occurs only for very rare, massive systems with $M_{500} \gta 10^{15} \msol$. To achieve the widest possible dynamic range, the second-order expressions could be interpolated using multiple pivot points, $M_{p,k}$, requiring values of $\betaone(M_{p,k}, z)$ and $\betatwo(M_{p,k}, z)$ to be provided. For light-cone applications, these derivatives could be modeled as low-order polynomials in redshift. The pivot masses we employ at the two demonstration redshifts correspond closely to those satisfying a fixed sky surface density condition, $dN(> \! M)/dz = {\rm const.}$, in a $\Lambda$CDM cosmology \citep{Evrard0110246, Mortonson1011.0004}. For cluster survey applications, this would seem a natural choice. While power-law scaling with log-normal covariance is supported for intrinsic properties of clusters, the available evidence is often limited. In particular, covariance among different signals is poorly understood \citep{Maughan1212.0858} and there are few constraints as to whether the slope and variance of a particular signal's scaling with mass is indeed constant with mass and redshift, as is assumed here \citep[\eg][provide a hint of evidence for curvature in the $\Lx$--$M$ relation]{BalagueraAntolnez1207.2138}. Redshift dependencies are easily incorporated into the existing framework by writing the slopes, $\boldalpha(z)$, intercepts, $\boldpi(z)$, and covariance, $C(z)$, as explicit functions of $z$. Extensions to the model that incorporate weak mass dependence in the scaling parameters are also possible. The model covariance can be interpreted as that among intrinsic properties of halos, but signals observed from real clusters inevitably include projection effects and noise, including potential bias, associated with signal detection and characterization. These effects, particularly projection for SZ and optical signatures, deserve further exploration \citep{NohCohn1204.1577, White1005.3022, Angulo1203.3216}. Using polynomial log-space approximations to the high-mass end of the cosmic mass function, we present analytic forms for statistics of multi-observable properties of the high-mass halos that host groups and clusters of galaxies. The model employs scaling laws between observables and mass that are power-law in the mean with log-normal covariance. The model provides quick estimation of cluster counts as a function of multiple observables and calculation of conditional likelihoods for observable-selected samples, both of which are directly relevant for joint survey analysis or follow-up observations. By comparing to a locally-convolved Tinker mass function, we show that the first-order model is generally accurate within a narrow range near the pivot mass, except for very high mass-scatter proxies. The second and third-order extensions provide increasingly wider coverage in observables irrespective of the mass scatter. The third-order model is nearly global in scope. The mass variance in a particular observable determines many expected features, as does the covariance between pairs of observables at fixed mass. As multi-wavelength surveys and dedicated follow-up campaigns provide increasingly rich, uniform samples of clusters, opportunities to apply this model to better constrain the statistical properties of massive halos will become apparent. Such knowledge will provide useful constraints on the physical processes that govern baryon evolution in massive halos.	14	3	1403.1456	The massive dark matter haloes that host groups and clusters of galaxies have observable properties that appear to be lognormally distributed about power-law mean scaling relations in halo mass. Coupling this assumption with either quadratic or cubic approximations to the mass function in log space, we derive closed-form expressions for the space density of haloes as a function of multiple observables as well as forms for the low-order moments of properties of observable-selected samples. Using a Tinker mass function in a Λ cold dark matter cosmology, we show that the cubic analytic model reproduces results obtained from direct, numerical convolution at the 10 per cent level or better over nearly the full range of observables covered by current observations and for redshifts extending to z = 1.5. The model provides an efficient framework for estimating effects arising from selection and covariance among observable properties in survey samples.	false	[ "observable properties", "survey samples", "halo mass", "properties", "multiple observables", "z", "current observations", "log space", "observables", "observable-selected samples", "forms", "power-law mean scaling relations", "redshifts", "host groups", "haloes", "a Λ cold dark matter cosmology", "covariance", "closed-form expressions", "galaxies", "selection" ]	11.462695	4.638595	-1
12205815	[ "Leon, Genly" ]	2011IJMPE..20...19L	[ "Phase Space of Anisotropic R<SUP>n</SUP> Cosmologies" ]	15	[ "Universidad Central de Las Villas, Department of Mathematics, Carretera a Camajuani km 5.5, Santa Clara, Villa Clara ZIP/Zone 54830, Cuba;" ]	[ "2015InJPh..89.1213B", "2016JCAP...11..051L", "2017GReGr..49...59A", "2017Prama..88...26B", "2017arXiv171007906B", "2018GReGr..50...24B", "2018PhRvD..98b4009C", "2019PhRvD..99b3520B", "2020CQGra..37s5014P", "2020EPJC...80..589P", "2020EPJC...80..816P", "2022GReGr..54...29B", "2022JCAP...09..042A", "2023AnPhy.45169245S", "2023PDU....4201355P" ]	[ "astronomy", "physics" ]	6	[ "Anisotropic Cosmology", "Phase Space", "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Extragalactic Astrophysics", "High Energy Physics - Theory" ]	[ "1930MNRAS..90..668E", "1967RvMP...39..862H", "1969CMaPh..12..108E", "1971CMaPh..23..137C", "1973ApJ...180..317C", "1983PhRvD..28.2118W", "1987NuPhB.292..784G", "1988CQGra...5..733D", "1990CQGra...7.1747B", "1991PhLB..254...44M", "1991PhLB..254...49C", "1993PhRvD..47.4315M", "1994MPLA....9.1897C", "1998PhRvD..57.6065B", "1999PhRvD..60b3502G", "1999PhRvD..60b4008N", "2002PhLB..524..177S", "2002PhRvD..65l6003S", "2003CQGra..20L.155B", "2003PhRvD..68h3512L", "2004PhRvL..92c1302K", "2005CQGra..22.4839C", "2006CQGra..23..001S", "2006CQGra..23.1913C", "2006CQGra..23.4915L", "2006PhRvD..74d3522K", "2007CQGra..24.5689G", "2007PhRvL..98g1301B", "2008ApJ...686..749K", "2008CQGra..25c5013G", "2008MNRAS.383..879A", "2008PhR...465...61T", "2008PhRvD..78d4011G", "2008arXiv0803.0905C", "2009ApJS..182..543A", "2009CQGra..26j5003G", "2009NuPhB.808..224S", "2010OAJ.....3...49C", "2010PhRvD..82h4015B", "2010RvMP...82..451S", "2011ApJS..192...14J", "2011CQGra..28c5010C", "2011CQGra..28f5008L" ]	[ "10.1142/S0218301311040037", "10.48550/arXiv.1403.1984" ]	1403	1403.1984_arXiv.txt	There exist a huge observational evidence that the expansion rate of our universe is now accelerating \cite{obs}. One way to explain this feature is to consider the extended gravitational theories known as $f(R)$-gravity (see \cite{Sotiriou:2008rp} and references therein). In such approach the Hilbert-Einstein action is generalized by replacing the Ricci scalar $R$ by functions of it. According to the observational evidence, the universe, highly inhomogeneous and anisotropic in earlier epochs, had evolved to the homogeneous and isotropic state we observe today with great accuracy. The robust approach to answer this question is to begin with an initially anisotropic universe and analyze if the universe evolves towards the future isotropization. Anisotropic but homogeneous cosmologies was known since a long time ago \cite{Misner:1974qy}. The most well-studied homogeneous but anisotropic geometries are the Bianchi type (see \cite{Ellis:1968vb} and references therein) and the Kantowski-Sachs metrics \cite{KS}, either in conventional or in higher-dimensional framework. For Bianchi I, Bianchi III, and Kantowski-Sachs geometries one can obtain a very good picture of homogeneous but anisotropic cosmology by using both numerical and analytical approaches and incorporating also the matter content (see \cite{Leon:2010pu} and references therein). In the present work we are interested in investigating the phase space of anisotropic $f(R)=R^n$ cosmologies focussing in Kantowski-Sachs geometries and modelling the matter content as a perfect fluid. The major emphasis is on the late-time stable solutions. We analyze the possibility of accelerating expansion at late times, and additionally, we determine conditions for the parameter $n$ for the existence of phantom behavior, contracting solutions as well as of cyclic cosmology. Furthermore, we analyze if the universe evolves towards the future isotropization. We make few comments about the feasibility to construct compact phase spaces for Bianchi III and Bianchi I background. We stress that the results of anisotropic $f(R)$ cosmology are expected to be different than the corresponding ones of $f(R)$-gravity in isotropic geometries, similarly to the differences between isotropic \cite{eddington} and anisotropic \cite{harrison} considerations in General Relativity. Additionally, the results are expected to be different from anisotropic General Relativity, too. As we see, anisotropic $f(R)$ cosmology can be consistent with observations.	As we saw, the universe at late times can result to a state of accelerating expansion, and additionally, for a particular $n$-range ($2<n<3$) it exhibits phantom behavior. Additionally, the universe has been isotropized, independently of the anisotropy degree of the initial conditions, and it asymptotically becomes flat. The fact that such features are in agreement with observations \cite{obs} is a significant advantage of the model. Moreover, in the case of radiation ($n=2$, $w=1/3$) the aforementioned stable solution corresponds to a de-Sitter expansion, and it can also describe the inflationary epoch of the universe. In our work we extracted our results without relying at all on the cosmic no-hair theorem, which is a significant advantage of the analysis \cite{Wald}. The Kantowski-Sachs anisotropic $R^n$-gravity can also lead to contracting solutions, either accelerating or decelerating, which are not globally stable. Thus, the universe can remain near these states for a long time, before the dynamics remove it towards the above expanding, accelerating, late-time attractors. The duration of these transient phases depends on the specific initial conditions. One of the most interesting behaviors is the possibility of the realization of the transition between expanding and contracting solutions during the evolution. That is, the scenario at hand can exhibit the cosmological bounce or turnaround. Additionally, there can also appear an eternal transition between expanding and contracting phases, that is we can obtain cyclic cosmology. These features can be of great significance for cosmology, since they are desirable in order for a model to be free of past or future singularities.	14	3	1403.1984	We construct general anisotropic cosmological scenarios governed by an f(R) = R<SUP>n</SUP> gravitational sector. Focusing then on some specific geometries, and modelling the matter content as a perfect fluid, we perform a phase-space analysis. We analyze the possibility of accelerating expansion at late times, and additionally, we determine conditions for the parameter n for the existence of phantom behavior, contracting solutions as well as of cyclic cosmology. Furthermore, we analyze if the universe evolves towards the future isotropization without relying on a cosmic no-hair theorem. Our results indicate that anisotropic geometries in modified gravitational frameworks present radically different cosmological behaviors compared to the simple isotropic scenarios.	false	[ "gravitational sector", "general anisotropic cosmological scenarios", "cyclic cosmology", ">n</SUP", "modified gravitational frameworks", "phantom behavior", "the simple isotropic scenarios", "radically different cosmological behaviors", "solutions", "late times", "conditions", "R<SUP", "a phase-space analysis", "expansion", "a perfect fluid", "an f(R", "that anisotropic geometries", "the matter content", "a cosmic no-hair theorem", "the future isotropization" ]	10.438164	0.676793	-1
513925	[ "Dudorov, A. E.", "Khaibrakhmanov, S. A." ]	2014Ap&SS.352..103D	[ "Fossil magnetic field of accretion disks of young stars" ]	35	[ "Chelyabinsk State University, Chelyabinsk, Russia", "Chelyabinsk State University, Chelyabinsk, Russia" ]	[ "2015AdSpR..55..843D", "2016A&AT...29..429D", "2016ARep...60..486P", "2016ApJ...833...92I", "2017A&A...602A..19L", "2017A&A...607A.104B", "2017ASPC..510...63A", "2017ASPC..510..114D", "2017MNRAS.464..586K", "2017PPNL...14..882K", "2018ApJ...869...45K", "2018Geosc...8..310W", "2018MNRAS.473.1427L", "2018RAA....18...90K", "2019EPJWC.20109004K", "2019MNRAS.487.5388D", "2020A&A...644A..74V", "2020ApJ...889...64A", "2020IAUS..345..295K", "2020IAUS..345..297D", "2020MNRAS.493.2101K", "2020MNRAS.497.1634G", "2020PSJ.....1...23K", "2021AJ....162....3C", "2021BLPI...48..312K", "2021Icar..35514127J", "2022ARep...66..872K", "2022MNRAS.516.4448K", "2022aems.conf..157K", "2023A&A...670A...6B", "2023A&A...678A.204S", "2023INASR...8..119B", "2023INASR...8..144K", "2024ARep...68...60S", "2024arXiv240114180K" ]	[ "astronomy" ]	5	[ "Accretion", "Accretion disks", "Diffusion", "MHD", "Stars: circumstellar matter", "ISM: evolution", "magnetic fields", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1954ooss.book.....A", "1961ApJ...134..270O", "1968ApJ...152..971S", "1972ApJ...173...87N", "1972epcf.book.....S", "1973A&A....24..337S", "1977Ap&SS..51..153W", "1978PASJ...30..671N", "1978ppim.book.....S", "1979cmft.book.....P", "1981PThPS..70...35H", "1982MNRAS.199..883B", "1987ARA&A..25...23S", "1987NInfo..63...68D", "1991ApJ...376..214B", "1991SvA....35..342D", "1993ApJ...406..122A", "1994ApJ...427..987B", "1994MNRAS.267..235L", "1995ARep...39..790D", "1995ApJ...439..752C", "1996ApJ...457..355G", "1996ApJ...459..653R", "1996ApJ...463..656S", "1996ApL&C..34..363A", "1997ApJ...480..344G", "1999PASA...16..225D", "2000ApJ...543..486S", "2000MNRAS.315..762S", "2002MNRAS.329...18F", "2003A&A...410..611S", "2003A&AT...22...11D", "2004ApJ...600..279C", "2004MNRAS.348..355K", "2005Natur.438..466D", "2006A&A...457..343F", "2006ApJ...650..956V", "2006Sci...313..812G", "2007ARA&A..45..565M", "2007Ap&SS.311...35W", "2007ApJ...665..535S", "2008ARep...52..790D", "2008ApJ...689..532T", "2009A&A...504..461F", "2009ApJ...690..407K", "2009ApJ...701..737B", "2010A&A...515A..70D", "2010ARep...54..776S", "2010apf..book.....A", "2011ARA&A..49...67W", "2011ApJ...729...83Y", "2011ApJ...736..144B", "2012MNRAS.420.3139M", "2012MNRAS.422.2737W", "2012MNRAS.424.2097G", "2013ApJ...764...65M", "2013ApJ...764...66S", "2013ApJ...769...76B", "2013ApJ...770..151C", "2013ApJ...772...96B", "2013ApJ...775...73S", "2013Sci...342.1069P", "2014ApJ...785..127O" ]	[ "10.1007/s10509-014-1900-4", "10.48550/arXiv.1403.5513" ]	1403	1403.5513_arXiv.txt	Observations indicate that stars born at present time in magnetized rotating cores of molecular clouds~\citep{shu87, dudorov91, dudorov95, mckee07}. Centrifugal and electromagnetic forces lead to the formation of a disk-like structure during the gravitational collapse of the molecular cloud cores. Evolution of accretion disks depends on efficiency of angular momentum removal. Turbulence, magnetic braking and outflows are the most important mechanisms of angular momentum transfer in accretion disks of young stars. Turbulence in accretion disks comes probably from magnetorotational instability \citep[MRI, ][]{balbus91}. Magnetic braking mechanism is based on the process of angular momentum transfer by torsional Alfven waves \citep{alfven_book}. Centrifugally driven winds arise when ordered magnetic field lines are inclined more than 30 degrees from vertical \citep{blandford82}. Efficiency of angular momentum transport depends on strength and geometry of the magnetic field. Numerical simulations indicate that initial magnetic flux of molecular clouds cores is partially conserved during the process of star formation \citep[e.g.,][]{dud87}. \cite{dudorov08} shown that the initially uniform magnetic field acquires hour-glass geometry during protostellar cloud collapse. Collapsing protostellar cloud with magnetic field evolves into flat structure according to numerical simulations of \cite{dudorov03}. They shown that the collapse of the clouds with strong magnetic field switches to the magnetostatic contraction into oblate self-graviting structures. These numerical simulations show that accretion disks of young stars should have fossil magnetic field. In other words, the magnetic flux of accretion disks is the relic of parent protostellar clouds magnetic flux. There are some observational confirmations of predictions of the fossil magnetic field theory \citep{dudorov95}. \cite{girart06} found hour-glass shaped geometry of the magnetic field in the low-mass protostellar system NGC 1333 IRAS 4A using the Submillimeter Array polarimetry observations at 877 $\mu$m dust continuum emission. \cite{chapman13} found some evidence of correlation between core magnetic field direction and protostellar disk symmetry axis in several low-mass protostellar cores using the SHARP polarimeter at Caltech Submillimeter Observatory at 350 $\mu$m dust continuum emission. \cite{donati05} probably detected magnetic field with azimuthal component of order of 1 kGs in the FU Orionis accretion disk. Observational data are scanty, so theoretical investigation of magnetic field of accretion disks of young stars is the important problem. Dragging of the magnetic field in the accretion disks has been investigated in several works. \cite{lubow94}, \cite{agapitou96} considered bending of the initially vertical weak magnetic field in infinitesimally thin viscous disk. They pointed out that inclination of poloidal magnetic fields lines and wind launching possibility depend on magnetic diffusion efficiency which is characterized by the magnetic Prandtl number $\mathcal{P}$. \cite{RRS96} showed that in the case of rotating disk, in addition to radial magnetic field component, strong toroidal component of the magnetic field is also generated when magnetic diffusion is feeble, $\mathcal{P}\gg 1$. \cite{shalybkov00} investigated steady-state structure of the accretion disk at the presence of dynamically important magnetic field. They found that sufficient inclination of the poloidal magnetic field lines to rotation axis is achieved also in the case of intermediate magnetic diffusion efficiency ($\mathcal{P}\simeq 1$). \cite{guilet12} pointed out that the accretion disk vertical structure has great effect on the magnetic flux evolution. In the limit of high conductivity, \cite{okuzumi13a} obtained analytical profile of the vertical component of the magnetic field, $B \propto r^{-2}$, that follows from magnetic flux conservation and gives upper estimation of the magnetic field strength in accretion disks, 0.1 Gs at 1 AU and $\sim$1 mGs at 10 AU. Ohmic dissipation (OD), magnetic ambipolar diffusion (MAD), turbulent diffusion and buoyancy are the basic dissipation mechanisms limiting fossil magnetic field during protostellar cloud collapse and subsequent accretion \citep[e.g.,][]{dud87}. Efficiency of OD and MAD depends on the ionization fraction. Cosmic rays \citep{gammie96} and stellar X-rays \citep{igea97} ionize effectively only surface layers of accretion disks of young stars. Regions of low ionization fraction (``dead'' zones) arise near accretion disk midplane under such circumstances. Magnetic diffusion suppresses MRI in the ``dead'' zones. Turbulence attenuation in the ``dead'' zones favours matter accumulation, gravitational instability and planet formation \citep[e.g.,][]{armitage_book}. There are three important non-ideal magnetohydrodynamical (MHD) effects operating in accretion disks: OD, MAD and Hall effect \citep[e.g.,][]{wardle07}. Efficiency of MAD and Hall effect depends on magnetic field strength. Geometry of the magnetic field also plays crucial role in the producing of MRI-induced turbulence. \cite{simon13b} showed that the vertical magnetic field enhances turbulent stresses comparing to the case with the purely toroidal magnetic field. Global numerical investigations of magnetized accretion disks and ``dead'' zones were performed in the ideal MHD limit \citep{fromang06} and in resistive limit \citep{dzyurkevich10} with magnetic field strength being free parameter. \cite{bai11}, \cite{simon13a, simon13b} performed calculations with MAD in the frame of the local shearing-box approximation with fixed magnetic field strength and/or geometry. Numerical modelling of magnetized accretion disks is still challenging problem, so semi-analytical models are widely used such as the minimum mass solar nebula model \citep[MMSN, e.g., ][]{weidenschiling77} and $\alpha$-model of \cite{ss73}. Investigations of ``dead'' zones in accretion disks in most cases use MMSN model \citep{sano00, bai09, mohanty13} or $\alpha$-model \citep{gammie96, fromang02, terquem08, martin12}. Large-scale magnetic field is ignored in these models. Prescribed magnetic field is also used in the models of magnetically driven winds \citep[e.g.,][]{CSD_book}. Recently, \cite{bai13a} and \cite{bai13b} performed numerical simulations in frame of local ``shearing-box'' approximation and showed that angular momentum transport is driven by magnetocentrifugal wind in the laminar ``dead'' zones, while MRI operates outside these regions. There are several papers concerning computation of the magnetic field of accretion disks. \cite{dekool99} constructed accretion disk model taking into account MHD turbulence. In order to incorporate magnetic field into the model, they used proportionality of the magnetic stresses and the Reynolds stresses \citep{stone96}. \cite{martin12} used similar approach in order to determine ``dead'' zones boundaries. \cite{vorobyov06} estimated vertical magnetic field $B_z$ in the accretion disk assuming flux freezing and using the mass-to-flux ratio $\lambda=2\pi G^{1/2}\Sigma / B_z$ \citep{nakano78}, where $\Sigma$ -- accretion disk surface density. \cite{shu07} made analytical estimations of vertical component of magnetic field from the equation of centrifugal balance in protoplanetary disk diluted by poloidal magnetic tension. They neglected azimuthal magnetic field component. We investigate the strength and geometry of the fossil large-scale magnetic field of accretion disks of young stars taking into account ohmic diffusion, magnetic ambipolar diffusion and buoyancy. The accretion disk model is based on the approximations of $\alpha$-model \citep{ss73}. We adopt kinematic approach and neglect magnetic field influence on the accretion disk structure. Ionization equations include shock ionization by X-rays, cosmic rays and radionuclides, thermal ionization, radiative recombinations and recombinations on the dust grains. Characteristics of ``dead'' zones in accretion disks of young stars are investigated in the frame of the elaborated model. The paper is organized as follows. In section \ref{Sec:Model}, we formulate basic equations of our model. In section \ref{Sec:AnalytResults}, we obtain and analyse analytical solution of the model equations. Results of numerical calculations of the magnetized accretion disks structure are presented in the section \ref{Sec:NumResults}. In section \ref{Sec:DZ}, results of investigation of ``dead'' zones characteristics are presented. We summarize and discuss our results in section \ref{Sec:Discussion}.	\label{Sec:Discussion} We elaborate kinematic MHD model of accretion disks of young stars. The model is based on Shakura and Sunyaev approximations. Magnetic field is calculated from stationary induction equation taking into account the ohmic diffusion (OD), magnetic ambipolar diffusion (MAD) and buoyancy. Ionization fraction is determined taking into consideration thermal ionization of metals and shock ionization by cosmic rays, X-rays and radioactive elements. Recombinations on dust grains, radiative recombinations and dust evaporation are included in the ionization model. Accretion disk sizes and masses calculated in the frame of our model agree with observations. Inner boundary of the accretion disk is placed at the distances $r=0.03-0.09$ AU (several stellar radii), outer boundary is placed at $r=70-270$ AU depending on stellar mass in the interval 0.5-2 $M_{\odot}$. Corresponding masses of accretion disks equal 0.003-0.17 of solar mass. We derive an analytical solution of the model equations in the case when dependence of ionization fraction on gas density is the power-law function. Attenuation of cosmic rays and X-rays leads to decrease of ionization fraction to extremely low values $\sim 10^{-16}-10^{-14}$ in the inner dense cold regions of accretion disk in the case of recombinations on the dust grains. Magnetic diffusion is efficient in this region of lowest ionization fraction (``dead'' zone). OD develops in the ``dead'' zone mainly where ionization fraction is minimal. MAD operates in the outer regions of the accretion disk. Magnetic field is frozen-in near the inner boundary of the accretion disk, where thermal ionization of metals with low ionization potential takes place. The fossil magnetic field of accretion disks of young stars has complex geometry. Initially uniform magnetic field of parent protostellar clouds transforms in accretion disk. All three components of the magnetic field are not zero. Differential rotation generates azimuthal component of the magnetic field. Accretion generates radial component of the magnetic field. We derive the following estimation for geometrically thin disk: $B_r\simeq 2/3\,\alpha B_{\varphi}$, were $\alpha$ is Shakura and Sunyaev parameter. Hence, radial component $B_r$ typically is much smaller than the azimuthal component $B_{\varphi}$ while their dependence on radial distance is the same. The magnetic field is qiasi-poloidal in the dusty ``dead'' zones. Estimations of the characteristic times of magnetic diffusion show that times of OD and MAD are less than lifetime of accretion disk in ``dead'' zones. Therefore, magnetic field geometry remains purely poloidal in this region. Strength of the vertical magnetic field is $B_z\simeq (10-20)$ mGs at $r=3$ AU depending on stellar mass, $M=0.5-2\,M_{\odot}$. Sedimentation of charged dust grains is possible in the ``dead'' zones along magnetic field lines. Magnetic field is coupled to the matter in the inner regions, $r\lesssim 0.3$ AU. Magnetic field is quazi-azimuthal in this region. Strength of magnetic field is comparable with stellar magnetic field strength at the accretion disk inner boundary, which equals 5-30 Gs depending on stellar mass. This means, that accretion disk inner boundary must be determined taking into account pressure of the accretion disk magnetic field. Interaction of the stellar magnetic field and accretion disk magnetic field near the accretion disk inner boundary may lead to current sheet formation that can become an additional source of X-ray activity of stellar magnetospheres. Magnetic field is quazi-azimuthal or quasi-radial in the outer regions of the accretion disk depending on intensity of ionization mechanisms. Growth of dust grains to $a_d\geq 1\,\mu$m, or increase of cosmic rays ionization rate to $10^{-16}\,\mathrm{s}^{-1}$, or increase of X-rays luminosity to $10^{32}\,\mathrm{erg}\,\mathrm{s}^{-1}$ are necessary to form quasi-radial magnetic field. The quasi-radial geometry of the magnetic field may be responsible for generation of magneto-centrifugal winds in the outer regions of the accretion disks owing to the criterion of \cite{blandford82}. Midplane ionization fraction does not fall below $10^{-11}$ in the case of radiative recombinations. Magnetic diffusion is inefficient and magnetic field is coupled to gas throughout the accretion disk. Vertical magnetic field is proportional to the surface density in this case. Our calculations show that $B_z(3\,\mbox{AU})=(0.03-0.13)$ Gs depending on stellar mass. Magnetic field is quasi-azimuthal, and $B_{\varphi}\simeq B_z$ at the distances less than several AU. Magnetic field is quasi-radial and magnetically-driven outflows are possible in the outer regions of accretion disks. Complexity of the fossil magnetic field geometry in accretion disks of young stars makes its investigation with Zeeman experiments and polarization measurements difficult. It requires resolution 0.1 AU at minimal distance 1 pc that corresponds to angular resolution $10^{-6}$ angular seconds. However, it is doubtful that, at present time, separation of Zeeman lines splitting from turbulent broadening of the lines in the accretion disk is possible. This is feasible for the ``dead'' zones only. ``Dead'' zones are very attractive for observations, both in terms of measurements of magnetic field strength, considering it is poloidal, and in terms of Earth type planet formation. Turbulence is weak and dust sedimentation is efficient here. Our calculations show that MAD determines the outer boundary of the ``dead'' zone for stars with $M>0.5\,M_{\odot}$. Outer boundary of the ``dead'' zone is placed at the distances 3-21 AU depending on the stellar mass in the case of recombinations on dust grains. In the case when dust is present in the disk, surface density of active layers is $\sim 10\,\mathrm{g}\,\mathrm{cm}^{-2}$ at $r\sim 1$ AU for typical parameters, i.e. ``dead'' zone occupies significant part of accretion disk thickness. Angular momentum transport by MRI-induced turbulence is inefficient under such circumstances. Tension of large-scale magnetic field lines may be responsible for the angular momentum transport in this region. Mass of solid material inside ``dead'' zones in accretion disks of stars with $M\geq 1\, M_{\odot}$ is more than 3 $M_{\oplus}$. This mass is sufficient for formation of several embryos of the Earth type planets in this region. Collisional formation and growth of planetesemals are the possible mechanisms of embryos formation \citep{safronov_book}. Our work is the step forward in the investigations of magnetized accretion disks of young stars comparing to existed semi-analytical investigations. In order to calculate magnetic field strength and geometry, we incorporate induction equations into the accretion disk model. Adopted kinematic approximation allows us investigate magnetic field structure with the help of quite simple semi-analytical equations of the model. We investigate ohmic diffusion and magnetic ambipolar diffusion in detail without paying attention to the Hall effect that is believed to influence the evolution of MHD turbulence in accretion disks of young stars \citep[e.g.,][]{wardle12}. Turbulence in the accretion disks also must influence magnetic field dissipation and generation. We intend to investigate the Hall effect and turbulent diffusion influence on the fossil magnetic field strength and geometry in accretion disks in our next paper. In order to improve elaborated model, we need to take into account dynamical influence of magnetic field on the accretion disk structure.	14	3	1403.5513	We elaborate the model of accretion disks of young stars with the fossil large-scale magnetic field in the frame of Shakura and Sunyaev approximation. Equations of the MHD model include Shakura and Sunyaev equations, induction equation and equations of ionization balance. Magnetic field is determined taking into account ohmic diffusion, magnetic ambipolar diffusion and buoyancy. Ionization fraction is calculated considering ionization by cosmic rays and X-rays, thermal ionization, radiative recombinations and recombinations on the dust grains. Analytical solution and numerical investigations show that the magnetic field is coupled to the gas in the case of radiative recombinations. Magnetic field is quasi-azimuthal close to accretion disk inner boundary and quasi-radial in the outer regions. Magnetic field is quasi-poloidal in the dusty "dead" zones with low ionization degree, where ohmic diffusion is efficient. Magnetic ambipolar diffusion reduces vertical magnetic field in 10 times comparing to the frozen-in field in this region. Magnetic field is quasi-azimuthal close to the outer boundary of accretion disks for standard ionization rates and dust grain size a <SUB>d</SUB>=0.1 μm. In the case of large dust grains ( a <SUB>d</SUB>>0.1 μm) or enhanced ionization rates, the magnetic field is quasi-radial in the outer regions. It is shown that the inner boundary of dusty "dead" zone is placed at r=(0.1-0.6) AU for accretion disks of stars with M=(0.5-2) M <SUB>⊙</SUB>. Outer boundary of "dead" zone is placed at r=(3-21) AU and it is determined by magnetic ambipolar diffusion. Mass of solid material in the "dead" zone is more than 3 M <SUB>⊕</SUB> for stars with M≥1 M <SUB>⊙</SUB>.	false	[ "Magnetic field", "vertical magnetic field", "enhanced ionization rates", "Magnetic ambipolar diffusion", "magnetic ambipolar diffusion", "standard ionization rates", "dust grain size", "large dust grains", "thermal ionization", "ionization balance", "low ionization degree", "ionization", "accretion disk", "accretion disks", "radiative recombinations", "ohmic diffusion", "account ohmic diffusion", "recombinations", "induction equation", "young stars" ]	10.183529	12.679779	-1
482793	[ "Chung, Sun-Ju", "Lee, Chung-Uk", "Koo, Jae-Rim" ]	2014ApJ...785..128C	[ "Detection of Planets in Extremely Weak Central Perturbation Microlensing Events via Next-generation Ground-based Surveys" ]	2	[ "Korea Astronomy and Space Science Institute 776, Daedeokdae-ro, Yuseong-Gu, Daejeon 305-348, Korea", "Korea Astronomy and Space Science Institute 776, Daedeokdae-ro, Yuseong-Gu, Daejeon 305-348, Korea", "Korea Astronomy and Space Science Institute 776, Daedeokdae-ro, Yuseong-Gu, Daejeon 305-348, Korea" ]	[ "2016ApJ...826...90C", "2018exha.book.....P" ]	[ "astronomy" ]	4	[ "gravitational lensing: micro", "planets and satellites: general", "Astrophysics - Earth and Planetary Astrophysics" ]	[ "1998ApJ...500...37G", "1999A&A...349..108D", "2002ApJ...572..521A", "2002MNRAS.335..159R", "2005ApJ...628L.109U", "2005ApJ...630..535C", "2006ApJ...644L..37G", "2006ApJ...650..432C", "2006Natur.439..437B", "2008ApJ...684..663B", "2008Sci...319..927G", "2009ApJ...698.1826D", "2009ApJ...705..386C", "2010ApJ...710.1641S", "2010ApJ...711..731J", "2010ApJ...720.1073G", "2010SPIE.7733E..3FK", "2011ApJ...728..120M", "2011ApJ...741...22M", "2012ApJ...754...73B", "2012ApJ...755..102Y", "2012ApJ...757..119B", "2013ApJ...762L..28H", "2013ApJ...768..129C", "2013ApJ...769...77Y", "2013ApJ...779...91F", "2014ApJ...780..123S", "2014ApJ...782...47P", "2014ApJ...782...48T" ]	[ "10.1088/0004-637X/785/2/128", "10.48550/arXiv.1403.1330" ]	1403	1403.1330_arXiv.txt	High-magnification events for which the background source star passes close to the host star are very sensitive for detection of planets \citep{griest98}. This is because the central caustic induced by a planet is formed near the host star and thus produces central perturbations near the peak of the lensing light curve. \citet{rattenbury02} studied planet detectability in high-magnification events and they showed that Earth-mass planets may be detected with 1 m class telescopes. Hence, current microlensing follow-up observations have focused on high-magnification events and have resulted so far in the detection of 14 planets out of 25 extrasolar planets detected by microlensing (Udalski et al. 2005; Gould et al. 2006; Gaudi et al. 2008; Bennett et al. 2008; Dong et al. 2009; Janczak et al. 2010; Miyake et al. 2011; Bachelet et al. 2012; Yee et al. 2012; Han et al. 2013; Choi et al. 2013; Suzuki et al. 2013). High-magnification events are sensitive to the diameter of the source star because the source star passes close to the central caustic. If the source diameter is bigger than the central caustic and thus the finite-source effect is strong, central perturbations induced by the central caustic are greatly washed out, thus making it difficult to realize the existence of planets. Events MOA-2007-BLG-400 (Subo et al. 2009), MOA-2008-BLG-310 (Janczak et al. 2010), and MOA-2010-BLG-311 (Yee et al. 2013) were high-magnification events with strong finite-source effects. All three events had complete coverage over the peak, but a secure planet detection occurred in only two events, MOA-2007-BLG-400 and MOA-2008-BLG-310. Even though the event MOA-2010-BLG-311 had complete coverage over the peak, central perturbations around the peak were extremely weak with a fractional deviation of $\lesssim 2\%$ so that it gave rise to a $\Delta \chi^2 \sim 80$ of the best-fit planetary lens model from the single lens model. \citet{yee13} reported that the planetary signal of $\Delta \chi^2 \sim 80$ is below the detection threshold range of $\Delta \chi^2 = 350-700$ suggested by \citet{gould10}, and thus it is difficult to claim a secure detection of the planet. This suggests that extremely weak central perturbations (hereafter EWCPs) of the deviations $\lesssim 2\%$ produce planetary signals below the detection threshold and obstruct a confident detection of planets. Current follow-up observations are intensively monitoring high-magnification events and their photometric error reaches $\sim 1\%$ at the peak, but it is not enough to get a confident detection of planets in EWCP events with deviations $\lesssim 2\%$, as shown in the event MOA-2010-BLG-311. For confident planet detection in EWCP events, it is necessary to have both high cadence monitoring and high photometric accuracy better than those of current follow-up observation systems. The next-generation ground-based observation project, KMTNet (Korea Microlensing Telescope Network), satisfies the conditions. The KMTNet will use a 1.6 m wide-field telescope at each of three southern sites, Chile, South Africa, and Australia, to perform a 24 hr continuous observation toward the Galactic bulge \citep{kim10}. Each telescope has a $18 {\rm K} \times 18 {\rm K}$ CCD camera that covers a field of view (FOV) of $2\arcdeg \times 2\arcdeg$ and it will observe four fields with a total FOV of $4\arcdeg \times 4\arcdeg$, in which each field will be monitored with an exposure time of about 2 minutes and a detector readout time of about 30 seconds giving a cadence of 10 minutes (Kim et al. 2010; Atwood et al. 2012). Hence, KMTNet has high potential for the detection of planetary signals in EWCP events. Here, we study how well planets in EWCP high-magnification events can be detected by the KMTNet. This paper is organized as follows. In \S\ 2, we briefly describe the properties of the central caustic induced by a planet. In \S\ 3, we estimate the probability of occurrence of EWCP events with deviations $\leq 2\%$ in high-magnification events and the efficiency of detecting planets in the EWCP events using KMTNet. In \S\ 4, we discuss the observational limitations and other potential studies of the KMTNet. We summarize the results and conclude in \S\ 5.	We have estimated the probability of occurring EWCP events of $\delta \leq 2\%$ in high-magnification events of $A_{\rm max} \geq 100$ and the efficiency of detecting planets in the EWCP events using the next-generation ground-based observation project, KMTNet. From this study, we found that the EWCP events occur with a frequency of $> 50\%$ in the case of $\lesssim 100\ M_{\rm E}$ planets with separations of $0.2\ {\rm AU} \lesssim d \lesssim 20\ {\rm AU}$. This implies that most of high-magnification events of $A_{\rm max} \geq 100$ are EWCP events of $\delta \leq 2\%$, and thus it is important to resolve the EWCP events for the detection of many different planetary systems. We found that for main-sequence and subgiant stars, $\gtrsim 1\ M_{\rm E}$ planets in EWCP events of $\delta \leq 2\%$ can be detected $> 50\%$ in a certain range that varies depending on the planet mass. However, it is difficult to detect planets in EWCP events of bright stars like giant stars, because it is easy for KMTNet to be saturated around the peak of the EWCP events with a constant exposure time. EWCP events are caused by close, intermediate, and wide planetary systems with low-mass planets and close and wide planetary systems with massive planets. Therefore, we expect that a much greater variety of planetary systems than those already detected, which are mostly intermediate planetary systems with $s \sim 1$ regardless of the planet mass, will be significantly detected in the near future. \begin{figure}[t] \epsscale{1.0} \plotone{fig1.eps} \caption{\label{fig:one} Probability of occurrence of EWCP events with $\delta \leq 2\%$ in high-magnification events of $A_{\rm max} \geq 100$ for three different source stars, main-sequence, subgiant, and giant stars, as a function of the projected star-planet separation in units of $\thetae$, $s$, and plant/star mass ratio, $q$. The physical separation, $d$, and planet mass in units of Earth mass, $m_{\rm p}$, are also presented. The radii of the three source stars in units of $\thetae$ are $\rho_{\star} = 0.0018, 0.0036,$ and $0.018$, respectively. In each panel, different shades of grey represent the areas with the probabilities of $\geq 10\%$, $\geq 40\%$, $\geq 80\%$, and $100\%$, respectively. The vertical dot line indicates the separation of $s = 1$. } \end{figure} \begin{figure}[t] \epsscale{1.0} \plotone{fig2.eps} \caption{\label{fig:three} Planet detection efficiency of EWCP events with $\delta \leq 2\%$ for three different source stars. Different shades of grey represent the areas with the efficiencies of $\geq 1\%, \geq 10\%, \geq 40\%,$ and $\geq 80\%$, respectively. The white solid and dashed lines represent the set of points where the probabilities of occurring events with $\delta \leq 1\%$ and $\delta \leq 0.5\%$ are both $80\%$. The white region marked as $``A"$ represents the region in which high-magnification events with $\delta > 2\%$ (not EWCP events) occur. } \end{figure} \begin{figure}[t] \epsscale{1.0} \plotone{fig3.eps} \caption{\label{fig:four} Detection efficiency as a function of the Einstein timescale for the planetary system of $s = 0.5$ and $q = 10^{-4}$. The vertical dashed line indicates a timescale of $\te = 10$ days. } \end{figure} \begin{figure}[t] \epsscale{1.0} \plotone{fig4.eps} \caption{\label{fig:four} Example light curve of a $I = 21.9$ star highly magnified by the planetary lens system of $s = 2.3$ and $q = 2.4\times 10^{-4}$. The black solid and grey dashed lines in the top panel are the light curves of the planetary lensing and best-fit single lensing events. } \end{figure} \begin{deluxetable}{lccccccc} \tablecaption{Probability of the occurrence of EWCP events.\label{tbl-one}} \tablewidth{0pt} \tablehead{ & \multicolumn{7}{c}{Probability ($\%$)} \\ \cline{3-7}\\ Planet mass && main-sequence && subgiant && giant } \startdata 300.0$M_{\rm E}$ && 35.9 && 40.1 && 61.3 \\ 100.0$M_{\rm E}$ && 51.6 && 56.2 && 78.9 \\ 10.0$M_{\rm E}$ && 79.7 && 85.7 && 96.2 \\ 5.0$M_{\rm E}$ && 85.7 && 91.5 && 96.5 \\ 1.0$M_{\rm E}$ && 95.3 && 96.9 && 97.0 \\ 0.5$M_{\rm E}$ && 97.1 && 97.5 && 97.2 \enddata \end{deluxetable} \begin{deluxetable}{lccccccccccccc} \tablecaption{Detection efficiency for a main-sequence star \label{tbl-two}} \tablewidth{0pt} \tablehead{ Planet mass && range (AU) &&& Detection Efficiency (DE) && Probability (P) && Ratio (DE/P) \\ && ($s < 1$) &&&\\ && ($s > 1$) &&& ($\%$) && ($\%$) && ($\%$) } \startdata 300.0$M_{\rm E}$ && $0.2 \lesssim d \lesssim 0.6 $ &&& 53.9 && 75.3 && 71.6 \\ && $6.6 \lesssim d \lesssim 17.8 $ &&& 54.4 && 75.4 && 72.1 \\\\ % 100.0$M_{\rm E}$ && $0.3 \lesssim d \lesssim 0.8 $ &&& 54.7 && 76.7 && 71.3 \\ && $4.6 \lesssim d \lesssim 12.4 $ &&& 54.7 && 75.8 && 72.2 \\\\ % 10.0$M_{\rm E}$ && $0.6 \lesssim d \lesssim 1.6 $ &&& 50.9 && 68.7 && 74.1 \\ && $2.3 \lesssim d \lesssim 5.7 $ &&& 50.5 && 68.5 && 73.7 \\\\ % 5.0$M_{\rm E}$ && $0.8 \lesssim d \lesssim 1.7 $ &&& 51.7 && 70.0 && 73.9 \\ && $2.1 \lesssim d \lesssim 4.6 $ &&& 51.3 && 69.5 && 73.8 \\\\ % 1.0$M_{\rm E}$ && $1.3 \lesssim d \lesssim 2.8 $ &&& 58.0 && 75.7 && 76.6 \\\\ % 0.5$M_{\rm E}$ && $1.5 \lesssim d \lesssim 2.4 $ &&& 57.4 && 85.0 && 67.5 \enddata \tablecomments{ The physical Einstein radius of the assumed lens system is $r_{\rm E} = 1.9\ \rm AU$. The range represents the region with the efficiency $\gtrsim 10 \%$, where it is rather wide, as shown in Figure 2. The range is divided into $s < 1$ and $s > 1$ only for the cases where the region with the efficiency $\gtrsim 10\%$ is clearly separated based on $s = 1$. } \end{deluxetable} \begin{deluxetable}{lccccccccccccc} \tablecaption{Detection efficiency for a subgiant star \label{tbl-three}} \tablewidth{0pt} \tablehead{ Planet mass && range (AU) &&& Detection Efficiency (DE) && Probability (P) && Ratio (DE/P) \\ && ($s < 1$) &&& \\ && ($s > 1$) &&& ($\%$) && ($\%$) && ($\%$) } \startdata 300.0$M_{\rm E}$ && $0.2 \lesssim d \lesssim 0.6 $ &&& 64.3 && 79.7 && 80.7 \\ && $6.4 \lesssim d \lesssim 15.5 $ &&& 64.6 && 80.1 && 80.6 \\\\ 100.0$M_{\rm E}$ && $0.3 \lesssim d \lesssim 0.8 $ &&& 65.8 && 80.9 && 81.3 \\ && $4.5 \lesssim d \lesssim 10.7 $ &&& 66.0 && 80.5 && 82.0 \\\\ 10.0$M_{\rm E}$ && $0.7 \lesssim d \lesssim 1.6 $ &&& 64.8 && 81.2 && 79.8 \\ && $2.3 \lesssim d \lesssim 4.9 $ &&& 64.9 && 81.7 && 79.4 \\\\ 5.0$M_{\rm E}$ && $0.9 \lesssim d \lesssim 1.8 $ &&& 62.6 && 79.1 && 79.1\\ && $2.0 \lesssim d \lesssim 3.9 $ &&& 62.4 && 79.0 && 79.0\\\\ 1.0$M_{\rm E}$ && $1.5 \lesssim d \lesssim 2.5 $ &&& 55.4 && 85.6 && 64.7\\\\ 0.5$M_{\rm E}$ && $1.7 \lesssim d \lesssim 2.1 $ &&& 39.0 && 90.1 && 43.3 \enddata \end{deluxetable} \begin{deluxetable}{lccccccccccccc} \tablecaption{Detection efficiency for a giant star \label{tbl-four}} \tablewidth{0pt} \tablehead{ Planet mass && range (AU) &&& Detection Efficiency (DE) && Probability (P) && Ratio (DE/P) \\ && ($s < 1$) &&& \\ && ($s > 1$) &&& ($\%$) && ($\%$) && ($\%$) } \startdata 300.0$M_{\rm E}$ && $0.4 \lesssim d \lesssim 0.8 $ &&& 68.2 && 87.8 && 77.7 \\ && $4.4 \lesssim d \lesssim 9.8 $ &&& 67.0 && 89.6 && 74.8 \\\\ 100.0$M_{\rm E}$ && $0.6 \lesssim d \lesssim 1.3 $ &&& 69.9 && 86.7 && 80.6 \\ && $2.9 \lesssim d \lesssim 6.5 $ &&& 71.0 && 89.1 && 79.7 \\\\ 10.0$M_{\rm E}$ && $1.3 \lesssim d \lesssim 2.8 $ &&& 57.6 && 81.7 && 70.5 \\\\ 5.0$M_{\rm E}$ && $1.6 \lesssim d \lesssim 2.2 $ &&& 66.9 && 83.5 && 80.1 \enddata \end{deluxetable}	14	3	1403.1330	Even though the recently discovered high-magnification event MOA-2010-BLG-311 had complete coverage over its peak, confident planet detection did not happen due to extremely weak central perturbations (EWCPs, fractional deviations of <~ 2%). For confident detection of planets in EWCP events, it is necessary to have both high cadence monitoring and high photometric accuracy better than those of current follow-up observation systems. The next-generation ground-based observation project, Korea Microlensing Telescope Network (KMTNet), satisfies these conditions. We estimate the probability of occurrence of EWCP events with fractional deviations of <=2% in high-magnification events and the efficiency of detecting planets in the EWCP events using the KMTNet. From this study, we find that the EWCP events occur with a frequency of >50% in the case of <~ 100 M <SUB>E</SUB> planets with separations of 0.2 AU <~ d <~ 20 AU. We find that for main-sequence and sub-giant source stars, >~ 1 M <SUB>E</SUB> planets in EWCP events with deviations <=2% can be detected with frequency >50% in a certain range that changes with the planet mass. However, it is difficult to detect planets in EWCP events of bright stars like giant stars because it is easy for KMTNet to be saturated around the peak of the events because of its constant exposure time. EWCP events are caused by close, intermediate, and wide planetary systems with low-mass planets and close and wide planetary systems with massive planets. Therefore, we expect that a much greater variety of planetary systems than those already detected, which are mostly intermediate planetary systems, regardless of the planet mass, will be significantly detected in the near future.	false	[ "planets", "massive planets", "EWCP events", "confident planet detection", "lt;~ d", "intermediate planetary systems", "planetary systems", "lt;~", "lt;=2%", "%", "fractional deviations", "EWCP", "low-mass planets", "high photometric accuracy", "giant stars", "Korea Microlensing Telescope Network", "KMTNet", "sub-giant source stars", "the planet mass", "close and wide planetary systems" ]	14.255866	12.263137	3
1619399	[ "Planck Collaboration", "Arnaud, M.", "Atrio-Barandela, F.", "Aumont, J.", "Baccigalupi, C.", "Banday, A. J.", "Barreiro, R. B.", "Battaner, E.", "Benabed, K.", "Benoit-Lévy, A.", "Bernard, J. -P.", "Bersanelli, M.", "Bielewicz, P.", "Bonaldi, A.", "Bond, J. R.", "Borrill, J.", "Bouchet, F. R.", "Buemi, C. S.", "Burigana, C.", "Cardoso, J. -F.", "Casassus, S.", "Catalano, A.", "Cerrigone, L.", "Chamballu, A.", "Chiang, H. C.", "Colombi, S.", "Colombo, L. P. L.", "Couchot, F.", "Crill, B. P.", "Curto, A.", "Cuttaia, F.", "Davies, R. D.", "Davis, R. J.", "de Bernardis, P.", "de Rosa, A.", "de Zotti, G.", "Delabrouille, J.", "Dickinson, C.", "Diego, J. M.", "Donzelli, S.", "Doré, O.", "Dupac, X.", "Enßlin, T. A.", "Eriksen, H. K.", "Finelli, F.", "Frailis, M.", "Franceschi, E.", "Galeotta, S.", "Ganga, K.", "Giard, M.", "González-Nuevo, J.", "Górski, K. M.", "Gregorio, A.", "Gruppuso, A.", "Hansen, F. K.", "Harrison, D. L.", "Hildebrandt, S. R.", "Hivon, E.", "Holmes, W. A.", "Hora, J. L.", "Hornstrup, A.", "Hovest, W.", "Huffenberger, K. M.", "Jaffe, T. R.", "Jones, W. C.", "Juvela, M.", "Keihänen, E.", "Keskitalo, R.", "Kisner, T. S.", "Knoche, J.", "Kunz, M.", "Kurki-Suonio, H.", "Lähteenmäki, A.", "Lamarre, J. -M.", "Lasenby, A.", "Lawrence, C. R.", "Leonardi, R.", "Leto, P.", "Lilje, P. B.", "Linden-Vørnle, M.", "López-Caniego, M.", "Macías-Pérez, J. F.", "Maffei, B.", "Maino, D.", "Mandolesi, N.", "Martin, P. G.", "Masi, S.", "Massardi, M.", "Matarrese, S.", "Mazzotta, P.", "Mendes, L.", "Mennella, A.", "Migliaccio, M.", "Miville-Deschênes, M. -A.", "Moneti, A.", "Montier, L.", "Morgante, G.", "Mortlock, D.", "Munshi, D.", "Murphy, J. A.", "Naselsky, P.", "Nati, F.", "Natoli, P.", "Noviello, F.", "Novikov, D.", "Novikov, I.", "Pagano, L.", "Pajot, F.", "Paladini, R.", "Paoletti, D.", "Peel, M.", "Perdereau, O.", "Perrotta, F.", "Piacentini, F.", "Piat, M.", "Pietrobon, D.", "Plaszczynski, S.", "Pointecouteau, E.", "Polenta, G.", "Popa, L.", "Pratt, G. W.", "Procopio, P.", "Prunet, S.", "Puget, J. -L.", "Rachen, J. P.", "Reinecke, M.", "Remazeilles, M.", "Ricciardi, S.", "Riller, T.", "Ristorcelli, I.", "Rocha, G.", "Rosset, C.", "Roudier, G.", "Rubiño-Martín, J. A.", "Rusholme, B.", "Sandri, M.", "Savini, G.", "Scott, D.", "Spencer, L. D.", "Stolyarov, V.", "Sutton, D.", "Suur-Uski, A. -S.", "Sygnet, J. -F.", "Tauber, J. A.", "Terenzi, L.", "Toffolatti, L.", "Tomasi, M.", "Trigilio, C.", "Tristram, M.", "Trombetti, T.", "Tucci, M.", "Umana, G.", "Valiviita, J.", "Van Tent, B.", "Vielva, P.", "Villa, F.", "Wade, L. A.", "Wandelt, B. D.", "Zacchei, A.", "Zijlstra, A.", "Zonca, A." ]	2015A&A...573A...6P	[ "Planck intermediate results. XVIII. The millimetre and sub-millimetre emission from planetary nebulae" ]	14	[ "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-" ]	[ "2015A&A...574A.134V", "2016A&A...588A.106W", "2017A&A...603A..67S", "2017MNRAS.468.3450C", "2018A&A...617A..85G", "2018A&A...619A..94P", "2018ApJ...858...22G", "2018NewAR..80....1D", "2019PhPl...26e2901G", "2020AJ....159..210N", "2021ApJ...919..121H", "2023RMxAA..59..279Q", "2024MNRAS.529.1579C", "2024PhRvD.109f3530C" ]	[ "astronomy" ]	46	[ "planetary nebulae: general", "radio continuum: ISM", "submillimeter: ISM", "Astrophysics - Astrophysics of Galaxies" ]	[ "1965ApJ...141..745M", "1965AuJPh..18..187S", "1965Natur.206..809D", "1966ApJ...144..657T", "1966SvA.....9..705K", "1967ApJ...148..429T", "1967ApJ...149..377H", "1967MNRAS.135..139D", "1968IAUS...34...87T", "1969A&A.....3..156R", "1969ApJ...155..469K", "1969ApL.....3...87T", "1969AuJPh..22..545L", "1970A&A.....8..171R", "1970A&AS....1..281C", "1970MNRAS.150..359T", "1971MNRAS.153..315H", "1972AuJPh..25...91A", "1973AJ.....78..875T", "1973MNRAS.161..313H", "1974ApJ...188..257T", "1974ApJ...192..165S", "1975A&A....38..183M", "1977ApJ...211..475T", "1977ApJ...217..425M", "1978ApJ...220...25E", "1980ApJ...238..585C", "1980ApJ...238..892M", "1981AJ.....86.1619U", "1981ApJ...247L..67K", "1982A&AS...50..209M", "1982MNRAS.199..141C", "1984AJ.....89..501T", "1984ApJ...276..544K", "1984ApJ...285...89D", "1984MNRAS.208..517G", "1986A&A...157..191G", "1986ApJS...61....1B", "1987A&A...171..178T", "1988A&A...197L..15M", "1988A&AS...75..317S", "1988ioch.rept.....H", "1991A&A...250..179Z", "1991ApJ...376..654G", "1991ApJS...75....1B", "1991ApJS...75.1011G", "1991MNRAS.251..330W", "1991Obs...111...72L", "1992A&A...261..567O", "1992MNRAS.258..257H", "1993ApJS...88..173K", "1994A&A...281..161A", "1994ApJS...90..179G", "1994MNRAS.271...75S", "1995ApJ...446..279B", "1996AJ....111.1945D", "1996ApJS..103..427G", "1997A&AS..124..259R", "1997ESASP.402..105H", "1997MNRAS.287..799I", "1997pdpn.book.....G", "1998AJ....115.1693C", "1998ApJ...495L..23C", "1998ApJS..117..361C", "1998PASP..110..761F", "1999A&A...349..243A", "1999astro.ph.10475I", "2000ApJ...532..384V", "2000MNRAS.314...99V", "2001A&A...367..949B", "2001A&A...378..843K", "2001Ap&SS.275..113S", "2001ApJ...546L.123C", "2001MNRAS.323..343L", "2002A&A...394...59D", "2002AJ....123..346S", "2002ApJ...574..179R", "2003ApJ...586..344W", "2003IAUS..209..281V", "2003yCat.5114....0E", "2004A&A...416..955S", "2004AJ....128.2339O", "2004ASPC..313...89S", "2004ApJ...602..960S", "2004ApJ...603..595F", "2004ApJ...603..599C", "2004ApJ...617.1142S", "2004ApJS..154...10F", "2004ApJS..154..296H", "2005A&A...431..523B", "2005ApJ...618..919G", "2005ApJS..157..302M", "2005MNRAS.360..963M", "2005MNRAS.362.1199C", "2006A&A...452..523K", "2006ApJ...644.1171S", "2006ApJ...652..426H", "2006MNRAS.369.1603G", "2006MNRAS.370.2047L", "2007A&A...461.1019G", "2007AJ....134.1679O", "2007AJ....134.2035N", "2007ApJS..171...61H", "2007MNRAS.382..299P", "2007MNRAS.382.1607C", "2007PASJ...59S.369M", "2008A&A...482..529U", "2008A&A...483..519S", "2008A&A...489...49V", "2008A&A...490..715P", "2008ApJ...681.1296Z", "2008ApJ...689..194S", "2008ApJS..175..277D", "2008IAUS..252..197K", "2008MNRAS.386.1404U", "2009A&A...498..463P", "2009A&A...507.1517M", "2009AJ....138..691C", "2009ApJS..180..283W", "2009MNRAS.399.1126P", "2010A&A...511A..53V", "2010A&A...517A..95P", "2010A&A...518L.137V", "2010A&A...520A...1T", "2010A&A...520A...3M", "2010A&A...520A...4B", "2010A&A...520A...8L", "2010A&A...520A...9L", "2010A&A...520A..13R", "2010AJ....140.1868W", "2010MNRAS.402.2403M", "2010MNRAS.409..881P", "2011A&A...536A...1P", "2011A&A...536A...2P", "2011A&A...536A...3M", "2011A&A...536A...4P", "2011A&A...536A...5Z", "2011A&A...536A...6P", "2011A&A...536A...8P", "2011A&A...536A...9P", "2011A&A...536A..10P", "2011A&A...536A..11P", "2011A&A...536A..12P", "2011A&A...536A..13P", "2011A&A...536A..14P", "2011A&A...536A..15P", "2011A&A...536A..16P", "2011A&A...536A..17P", "2011A&A...536A..18P", "2011A&A...536A..19P", "2011A&A...536A..20P", "2011A&A...536A..21P", "2011A&A...536A..22P", "2011A&A...536A..23P", "2011A&A...536A..24P", "2011A&A...536A..25P", "2011A&A...536A..26P", "2011AJ....142...91R", "2011ApJ...738..174M", "2011ApJS..192...15G", "2011MNRAS.411.1245R", "2011MNRAS.416..790A", "2011RMxAA..47...83P", "2012ApJ...755...53Z", "2013A&A...556A..35T", "2013MNRAS.435.3462M", "2014A&A...571A...1P", "2014A&A...571A...2P", "2014A&A...571A...6P", "2014A&A...571A..28P" ]	[ "10.1051/0004-6361/201423836", "10.48550/arXiv.1403.4723" ]	1403	1403.4723_arXiv.txt	\label{sec:intro} The final phases of low to intermediate mass stars are characterized by periods of high mass loss that lead to the formation of dense circumstellar envelopes (CSEs), where physical conditions are ideal for dust to condense (during the Asymptotic Giant Branch, or AGB phase). Such envelopes can be very massive and, in some cases, the central object can be completely optically obscured. Eventually the mass loss stops and the central star becomes visible as the dusty shell disperses (this is the proto-planetary nebula, or PPN phase). During the subsequent evolutionary phases, the central star moves toward higher temperatures and, once the stellar temperature is high enough to ionise the surrounding medium, the object becomes a planetary nebula (PN) \citep{Kwok2008}. PNe are usually surrounded by a dusty envelope that is a remnant of the previous evolutionary phases and is partly ionised by the UV radiation from the central star. The material surrounding the central object consequently has quite a complex distribution, consisting of concentric shells. In these shells, the level of ionisation decreases with the distance to the central star: highly ionised species are close to the star, while the outer part of the nebula is characterized by molecules and dust. This characteristic circumstellar environment implies the presence of two important components in the spectral energy distribution (SED) of a typical PN, whose major contributions fall in the spectral range from the far-IR to the radio region. The ionised fraction of the CSE can be traced by its free-free emission, which makes PNe bright Galactic radio sources, with some of them reaching flux densities up to a few Jy. Dust thermally re-radiates the absorbed stellar light, showing a clear signature in the far-infrared (far-IR) spectrum, i.e., an IR excess in the colour-colour diagram created with data from the Infrared Astronomical Satellite (IRAS). Such a contribution is typically of the order of 40\,\% of the total flux from a PN \citep{Zhang1991}, and is larger in young PNe, since in more evolved PNe the circumstellar material has already dispersed. PNe and their progenitors are considered to be among the major sources of recycled material into the ISM and for this reason, they are regarded as key objects for studying the chemical evolution of the Galaxy. Before being released into the ISM, significant processing of the material contained in the PN envelopes is expected. Gas and dust are exposed to a very harsh environment: shocks will occur when the fast outflows developing at the beginning of the PN phase overtake the slow expanding AGB envelope, and the UV radiation field radiated by the central star could be very intense, with consequences on the ionisation/recombination equilibrium of the gas and on the dusty/molecular content of the envelope \citep{Hora2009}. It is therefore very important to establish not only how much of this processed material is returned to the ISM after the nebula disperses, but also its general properties and dominant chemistry. The aim of this work is to derive the physical characteristics of a sample of Galactic PNe, taking advantage of the unique frequency coverage provided by the \Planck\footnote{\Planck\ is a project of the European Space Agency -- ESA -- with instruments provided by two scientific Consortia funded by ESA member states (in particular the lead countries: France and Italy) with contributions from NASA (USA), and telescope reflectors provided in a collaboration between ESA and a scientific Consortium led and funded by Denmark.} measurements, which trace both the ionised and the dust components. We model the SEDs with particular attention to the continuum from the mid-IR to the radio. Results from such modelling will be used to derive important parameters of PN envelopes, i.e., the total mass, the ionised fraction and the properties of the dust component. In some cases, hints on the extended morphology of PN envelopes can be derived from a direct inspection of the \Planck\ maps. As a starting point, we have compiled a master catalogue of Galactic PNe for which pre-existing 30\GHz\ and/or 43\GHz\ measurements are available. Our catalogue, based on single-dish measurements from the Torun (30\GHz, \citealp{Pazderska2009}) and Noto (43\GHz, \citealp{Umana2008a}) surveys, consists of 119 PNe, and covers a large parameter space in terms of location with respect to the Galactic plane, distance, and evolutionary stage. Typical numbers for the surveys are a FWHM for the beam of 72\arcs\ and an rms noise of 5\,mJy for the One Centimetre Receiver Array-prototype (OCRA-p) observations, and a FWHM for the beam of 52\arcs\ and rms noise of 70\,mJy for the Noto telescope survey. This paper is organized as follows. Observations, consisting of \Planck\ and ancillary data, are presented in Sect. \ref{sec:obs}, while the methods to extract fluxes and results are described in Sect. \ref{sec:sed}. The adopted SED modelling and derived physical properties of the detected targets are illustrated in Sect. \ref{sec:properties}, while Sects. \ref{sec:crl618} and \ref{sec:helix} focus on two targets, namely CLR\,618 and NGC\,7293 (the Helix Nebula, Helix hereafter), whose characterization appear to be particularly interesting. Section \ref{sec:conclusions} concludes. \begin{table} \begingroup \newdimen\tblskip \tblskip=5pt \caption{Non-colour-corrected flux densities (Jy) from Planck Catalog of Compact Sources \citep[PCCS;][]{planck2013-p05}. } \label{tab:fluxes} \nointerlineskip \vskip -3mm \footnotesize \setbox\tablebox=\vbox{ \newdimen\digitwidth \setbox0=\hbox{\rm 0} \digitwidth=\wd0 \catcode`=\active \def{\kern\digitwidth} \newdimen\signwidth \setbox0=\hbox{+} \signwidth=\wd0 \catcode`!=\active \def!{\kern\signwidth} \halign{\hbox to 1in{#\leaderfil}\tabskip .35em& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& \hfil#\hfil& #\hfil\tabskip 0pt\cr \noalign{\doubleline\vskip 2pt} \omit Source \hfil& Coordinates& \multispan9\hfil Frequency (GHz) \hfil& Ancillary Data\cr \noalign{\vskip 4pt\hrule\vskip 6pt} \omit& Galactic& 28.4& 44.1& 70.4& 100& 143& 217& 353& 545& 857& Refs.\cr \noalign{\vskip 4pt\hrule\vskip 6pt} \omit NGC\,6369$^{\phantom \dag}_{\phantom \dag}$& \s002.432+05.847& \s1.9\phantom{0}$\pm$0.2\phantom{0}& \s1.7$\pm$0.3\phantom{0}& \s1.27$\pm$0.15& \s1.5\phantom{0}$\pm$0.1\phantom{0}& \s1.27$\pm$0.06& \s1.35$\pm$0.09& \s\phantom{0}2.4\phantom{0}$\pm$0.3\phantom{0}& \s\phantom{0}7.0\phantom{0}$\pm$0.9\phantom{0}& \dots& \s$^{3,~8,~12,~15,~16,~17,~19,~21,~22,~29,~30,~34,~40,~45,~47}_{50,~54}$ \cr \omit NGC\,6572$^{\phantom \dag}_{\phantom \dag}$& \s034.623+11.848& \s1.1\phantom{0}$\pm$0.1\phantom{0}& \dots& \s1.2\phantom{0}$\pm$0.2\phantom{0}& \s0.96$\pm$0.06& \s0.92$\pm$0.04& \s0.81$\pm$0.04& \s0.81$\pm$0.09& \dots& \dots& \s{$^{1,~2,~3,~4,~5,~6,~7,~8,~9,~10,~13,~16,~17,~18,~19,~21,~22,~23}_{26,~29,~30,~32,~33,~34,~35,~38,~39,~48,~50,~54,~55,~56}$}\cr \omit NGC\,7293$^{\mathrm a}$ $^{\phantom \dag}_{\phantom \dag}$& \s036.161$-$57.118& \s0.8\phantom{0}$\pm$0.1\phantom{0}& \s1.0$\pm$0.2\phantom{0}& \s0.88$\pm$0.15& \s1.6\phantom{0}$\pm$0.5\phantom{0}& \s1.6\phantom{0}$\pm$0.5\phantom{0}& \s3.4\phantom{0}$\pm$0.6\phantom{0}& \s10\phantom{.00}$\pm$1\phantom{.00}& \s26\phantom{.00}$\pm$3\phantom{.00}& \s73\phantom{.0}$\pm$9\phantom{.0}& \s{$^{2,~3,~6,~7,~16,~17,~18,~19,~23,~29,~52}$}\cr \omit NGC\,7009$^{\phantom \dag}_{\phantom \dag}$& \s037.762$-$34.571& \s0.6\phantom{0}$\pm$0.1\phantom{0}& \dots& \s0.8\phantom{0}$\pm$0.2\phantom{0}& \s0.54$\pm$0.06& \s0.40$\pm$0.04& \s0.42$\pm$0.04& \dots& \dots& \dots& \s{$^{2,~3,~6,~8,~9,~16,~17,~21,~26,~29,~30,~41,~45,~47,~50,~51}_{54,~55,~57,~59}$}\cr \omit NGC\,6853$^{\phantom \dag}_{\phantom \dag}$& \s060.836$-$03.696& \dots& \dots& \s0.97$\pm$0.15& \s0.87$\pm$0.07& \s0.65$\pm$0.05& \dots& \dots& \dots& \dots& \s{$^{1,~2,~3,~4,~5,~6,~7,~8,~9,~10,~13,~16,~17,~18,~20,~23,~38}_{50}$}\cr \omit NGC\,6720$^{\phantom \dag}_{\phantom \dag}$& \s063.170+13.978& \dots& \dots& \dots& \dots& \s0.31$\pm$0.04& \s0.33$\pm$0.03& \s\phantom{0}0.64$\pm$0.06& \s\phantom{0}1.6\phantom{0}$\pm$0.1\phantom{0}& \s\phantom{0}4.7$\pm$0.2& \s{$^{1,~2,~5,~6,~7,~8,~9,~10,~14,~16,~18,~20,~21,~32,~33,~38,~39}_{48,~50,~55,~60}$}\cr \omit NGC\,6826$^{\phantom \dag}_{\phantom \dag}$& \s083.568+12.792& \s0.35$\pm$0.08& \dots& \dots& \s0.32$\pm$0.04& \s0.25$\pm$0.03& \s0.26$\pm$0.02& \dots& \dots& \dots& \s{$^{6,~8,~10,~12,~16,~17,~21,~33,~38,~39,~48,~50,~55,~60}$}\cr \omit NGC\,7027$^{\phantom \dag}_{\phantom \dag}$& \s084.930$-$03.496& \s5.0\phantom{0}$\pm$0.9\phantom{0}& \s4.7$\pm$0.6\phantom{0}& \s4.9\phantom{0}$\pm$0.2\phantom{0}& \s4.4\phantom{0}$\pm$0.1\phantom{0}& \s4.25$\pm$0.06& \s4.3\phantom{0}$\pm$0.1\phantom{0}& \s\phantom{0}5.4\phantom{0}$\pm$0.3\phantom{0}& \s\phantom{0}6.3\phantom{0}$\pm$0.7\phantom{0}& \dots& \s{$^{13,~24,~25,~27,~28,~31,~36,~42,~43,~44,~46,~56}$}\cr \omit NGC\,6543$^{\phantom \dag}_{\phantom \dag}$& \s096.468+29.954& \s0.74$\pm$0.08& \s0.9$\pm$0.15& \s0.67$\pm$0.09& \s0.63$\pm$0.04& \s0.52$\pm$0.02& \s0.54$\pm$0.02& \s\phantom{0}0.59$\pm$0.04& \s\phantom{0}0.84$\pm$0.07& \s\phantom{0}2.6$\pm$0.2& \s{$^{2,~4,~6,~7,~8,~16,~17,~21,~22,~27,~38,~39,~47,~48,~49,~50}_{53,~55,~60}$}\cr \omit NGC\,40$^{\phantom \dag}_{\phantom \dag}$& \s120.016+09.868& \dots& \dots& \s0.78$\pm$0.15& \s0.45$\pm$0.06& \s0.36$\pm$0.04& \s0.30$\pm$0.04& \dots& \dots& \dots& \s{$^{6,~10,~12,~16,~17,~38,~39,~48,~50,~60}$}\cr \omit CRL\,618$^{\phantom \dag}_{\phantom \dag}$& \s166.446$-$06.527& \s0.7\phantom{0}$\pm$0.2\phantom{0}& \s1.4$\pm$0.2\phantom{0}& \s1.9\phantom{0}$\pm$0.2\phantom{0}& \s2.40$\pm$0.08& \s2.67$\pm$0.05& \s3.00$\pm$0.07& \s\phantom{0}4.9\phantom{0}$\pm$0.2\phantom{0}& \s10.4$\pm$0.3\phantom{0}& \s24.3$\pm$0.8& \s{$^{32,~56}$}\cr \omit IC\,418$^{\phantom \dag}_{\phantom \dag}$& \s215.212$-$24.283& \s1.6\phantom{0}$\pm$0.1\phantom{0}& \s1.5$\pm$0.2\phantom{0}& \s1.4\phantom{0}$\pm$0.2\phantom{0}& \s1.18$\pm$0.07& \s1.02$\pm$0.04& \s0.94$\pm$0.03& \s\phantom{0}0.84$\pm$0.08& \dots& \dots& \s{$^{2,~3,~4,~6,~7,~8,~9,~10,~11,~13,~15,~16,~17,~19,~21,~22,~30}_{35,~37,~42,~45,~50,~54,~55,~56,~58,~59}$}\cr \omit NGC\,3242$^{\phantom \dag}_{\phantom \dag}$& \s261.051+32.050& \s0.6\phantom{0}$\pm$0.1\phantom{0}& \dots& \dots& \s0.51$\pm$0.06& \s0.48$\pm$0.04& \s0.37$\pm$0.03& \s\phantom{0}0.44$\pm$0.06& \dots& \dots& \s{$^{2,~3,~4,~6,~8,~9,~10,~11,~12,~13,~16,~17,~21,~22,~29,~35}_{37,~45,~50,~51,~54,~55,~56,~59}$}\cr \noalign{\vskip 3pt\hrule\vskip 4pt} }} \endPlancktablewide \vspace{-2mm} \begin{list}{}{} \item[$^{\mathrm a}$] Helix Nebula \item[References:] (1) \citealp{davies_etal65}; (2) \citealp{menon_terzian65}; (3) \citealp{slee_orchiston65}; (4) \citealp{khromov_66}; (5) \citealp{terzian66}; (6) \citealp{davies_etal67}; (7) \citealp{hughes_67}; (8) \citealp{thompson_etal67}; (9) \citealp{terzian68}; (10) \citealp{kaftan-kassim_69}; (11) \citealp{le_marne_69}; (12) \citealp{ribes69}; (13) \citealp{terzian69}; (14) \citealp{colla_etal70}; (15) \citealp{rubin_70}; (16) \citealp{thomasson_davies70}; (17) \citealp{higgs_71}; (18) \citealp{aller_milne72}; (19) \citealp{Higgs1973}; (20) \citealp{terzian_dickey73}; (21) \citealp{sistla_etal74}; (22) \citealp{terzian_etal74}; (23) \citealp{milne_aller75}; (24) \citealp{Telesco1977}; (25) \citealp{Elias1978}; (26) \citealp{cohen_barlow1980}; (27) \citealp{Moseley1980}; (28) \citealp{Ulich1981}; (29) \citealp{calabretta82}; (30) \citealp{milne_aller82}; (31) \citealp{Gee1984}; (32) \citealp{turner_terzian84}; (33) \citealp{bennet_etal86}; (34) \citealp{gathier_etal86}; (35) \citealp{taylor_etal87}; (36) \citealp{steppe_etal88}; (37) \citealp{wright_otrupcek90}; (38) \citealp{becker_etal91}; (39) \citealp{gregory_condon1991}; (40) \citealp{large_etal91}; (41) \citealp{wright_etal91}; (42) \citealp{Hoare1992}; (43) \citealp{knapp_etal93}; (44) \citealp{altenhoff_etal94}; (45) \citealp{griffith_etal94}; (46) \citealp{sandelll94}; (47) \citealp{douglas_etal96}; (48) \citealp{gregory_etal1996}; (49) \citealp{Rengelink_etal97}; (50) \citealp{condon_kaplan98}; (51) \citealp{De_Breuck_etal02}; (52) \citealp{Casassus2004}; (53) \citealp{Klaas2006}; (54) \citealp{casassus_etal07}; (55) \citealp{healey_etal07}; (56) \citealp{DiFrancesco2008}; (57) \citealp{vollmer_etal08}; (58) \citealp{wright_etal09}; (59) \citealp{murphy_etal10}; (60) \citealp{vollmer_etal10}. \end{list} \endgroup \end{table}	\label{sec:conclusions} The remarkable frequency coverage of \Planck\ has allowed us to collect important data for a small sample of Galactic PNe. These PNe, being among the brightest of such sources, are the best studied in our Galaxy. Nevertheless, a comprehensive picture of their CSEs in both their main components, i.e., ionised gas and dust, has been provided for the first time. In particular, the evaluation of the emission from ionised gas (free-free) allows us to constrain the thermal dust emission. The modelling of the SEDs provides us with good estimates of the physical parameters of the CSEs. From our modelling, the density, spatial distribution and total mass of ionised gas have been derived, as well as the internal radius, the extent of the dusty envelope and its mass content. One interesting result is that, in general in the studied targets, dust and ionised gas appear to be partially co-spatial, which implies the existence of some kind of shielding mechanism to allow the dust to form, or at least to survive, in the harsh environment created by the strong UV radiation field of the central object. Such shielding can be provided by a dusty disk or a molecular torus, as pointed out by interferometric observations of young, bi-polar PNe, i.e., CRL\,618, or can be in the form of clumps of material, as observed in older objects such as the Helix or NGC\,6720. These shielding structures must be very common and may well be debris from material existing before the nebula was formed. Thanks to the multi-frequency \Planck\ measurements, the SED of the variable source CRL\,618 was modelled for the first time, including both ionised gas and thermal dust contributions. This allowed us to derive important physical parameters of the CSE without variability effects. Finally, in the case of the Helix, \Planck\ maps enable us to perform a morphological study of the extended circumstellar material associated with the evolved PN. The dust emission was fully mapped for the first time and three main components were found, including an extended structure already seen in $H_\alpha$ observations, probably related to a region where the slow expanding envelope interacts with the surrounding ISM. A comparison was also performed between the morphology of the dust component as traced by \Planck, of the molecular gas traced by H$_{2}$ near-IR observations, and of the ionised gas traced by radio (1.4\GHz) observations. While the dust and H$_{2}$ share a comparable morphology, the ionised gas appears more concentrated in the inner ring of the nebula. Now that Planck has robustly determined the large-scale emission of this PN sample over a factor of 30 in wavelength, there is a firm basis within which to carry out detailed follow-up studies of the higher spatial and spectral-resolution properties. In particular, the results presented in this paper will provide us with a useful framework within which to plan future projects with ALMA, focusing on detailed morphological studies of both the ionised and dusty components. \begin{table*}[tmb] \appendix	14	3	1403.4723	Late stages of stellar evolution are characterized by copious mass-loss events whose signature is the formation of circumstellar envelopes (CSE). Planck multi-frequency measurements have provided relevant information on a sample of Galactic planetary nebulae (PNe) in the important and relatively unexplored observational band between 30 and 857 GHz. Planck enables the assembly of comprehensive PNe spectral energy distributions (SEDs) from radio to far-IR frequencies. Modelling the derived SEDs provides us with information on physical properties of CSEs and the mass content of both main components: ionized gas, traced by the free-free emission at cm-mm waves; and thermal dust, traced by the millimetre and far-IR emission. In particular, the amount of ionized gas and dust has been derived here. Such quantities have also been estimated for the very young PN CRL 618, where the strong variability observed in its radio and millimetre emission has previously prevented constructing its SED. A morphological study of the Helix Nebula was also performed. Planck maps reveal, for the first time, the spatial distribution of the dust inside the envelope, allowing us to identify different components, the most interesting of which is a very extended component (up to 1 pc) that may be related to a region where the slow expanding envelope is interacting with the surrounding interstellar medium.	false	[ "different components", "circumstellar envelopes", "thermal dust", "SED", "dust", "ionized gas", "comprehensive PNe spectral energy distributions", "relevant information", "radio", "the slow expanding envelope", "the surrounding interstellar medium", "physical properties", "information", "CSE", "SEDs", "Planck multi-frequency measurements", "cm-mm waves", "both main components", "PNe", "its radio and millimetre emission" ]	10.283957	10.119009	-1
525804	[ "Muscat, Daniel" ]	2014arXiv1403.4209M	[ "High-Performance Image Synthesis for Radio Interferometry" ]	7	[ "-" ]	[ "2018PASP..130a4503M", "2019AcASn..60...12L", "2020bda..book..271Z", "2021arXiv211104141M", "2022RAA....22j5010Y", "2023A&C....4500767M", "2023RASTI...2...91G" ]	[ "astronomy" ]	11	[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]	[ "1920ApJ....51..257M", "1921ApJ....53..249M", "1958PIRE...46...97B", "1962Natur.194..517R", "1971PhDT.......153B", "1974A&AS...15..417H", "1974AJ.....79...11T", "1980A&A....84...75F", "1980A&A....89..377C", "1987RvGeo..25..867R", "1992A&A...261..353C", "1995PhDT.......238B", "2006A&A...446..747G", "2008ISTSP...2..647C", "2010A&A...524A..42P", "2010CoPhC.181.1707E", "2010PASP..122.1353O", "2012SPIE.8500E..0LC", "2013A&A...556A...2V", "2013ApJ...770...91B", "2013MNRAS.430.2390B", "2013PASA...30...31B" ]	[ "10.48550/arXiv.1403.4209" ]	1403	1403.4209_arXiv.txt	The first interferometer used in astronomy dates back to 1921 when Albert Abraham Michelson together with Francis G. Pease made the first diameter measurements of the star Betelgeuse \cite{Michelson1921}. The setting used, known as the \textit{Michelson Interferometer}, is today the basis of modern interferometers. Michelson had been discussing the use of interferometry in astronomy for at least 30 years before the experiment was done \cite{CI_Michelson_stellar,Michelson1920}. Throughout the century, the technology of radio interferometers has made enormous advancement and up to this day extraordinarily ambitious projects have been commissioned, and more are on the pipeline. An example of a recently commissioned radio interferometer is the \textit{LOw-Frequency ARray} (LOFAR) \cite{Haarlem2013}. The \textit{Square Kilometre Array} (SKA) \cite{2013,Dewdney2013,P.E.Dewdney2011,P.E.Dewdney2010,R.T.Schilizzi2007} is the most ambitious project currently under-way. Radio interferometers can be used for various purposes, even for applications outside the scope of astronomy, such as geophysics \cite{Carter1988,Robertson1987}. The scope of this thesis is limited to the use of the interferometer as a measuring device for the intensity distribution of the sky over the celestial sphere. The measured quantity of the interferometer is known as \textit{visibility}, and the thesis' main focus is on how to recover the said intensity distribution from such measured visibility data. An interferometer is made up of an array of N$\geq$ 2 antennas, and differently from a single dish antenna it achieves sub-arcsecond resolutions with high accuracy. The maximum angular resolution of a single dish $\theta_m$ is limited by its diffraction limit. The limit is inversely proportional to the diameter $D$ of the dish, that is $\theta_m\propto 1/D$. On the other hand, the angular resolution of the interferometer is limited by the distance between the furthest two antennas in the array $B_{max}$ in the same inverse proportional way, that is $\theta_m\propto 1/B_{max}$. Achieving sub-arcsecond resolution in single-dish antennas requires large diameters that are prohibitory. For interferometers, achieving sub-arcsecond resolution is just a matter of having antennas as far away as possible from each other. If necessary, part of the array can be orbiting in space such as in the case of the \textit{Very Long Baseline Interferometry} (VLBI) \textit{Space Observatory Programme} (VSOP) \cite{Hirabayashi1998}. The basic measurement device (based on the \textit{Michelson interferometer}) is composed of just two elements. Each possible two element combination of N $\geq$ 2 antennas forms a two element independent measuring device. During an observation, the antenna array tracks a position on the celestial sphere normally within the field of view of the observation. Each basic element makes a measurement of the intensity distribution of the sky in the form of visibility values. $N(N-1)/2$ visibility readings are done simultaneously by the whole $N$-element antenna array and each reading differs in the geographical set-up of the basic device. The geographical set-up is described by the \textit{baseline} which is the vector covering the distance between the two antennas. Earth rotation changes the directions of the baseline (assuming a frame of reference stationary with the celestial sphere), and this is taken advantage of by taking subsequent visibility readings \cite{Ryle1962}. As it will be shown later on in this chapter, a Fourier relationship exists between visibility as a function of baseline and the intensity distribution, provided that certain conditions are met. To take advantage of such a relationship is not an easy computational task and is today an active area of research which this thesis is part of. This introduction aims to give a brief on the theory of interferometry, defines the measurement equation of the interferometer and discusses an image synthesis pipeline commonly used. The brief serves as a preamble for the discussion on motivations, aims and objectives of this thesis, which is done in the penultimate section. The chapter is concluded by giving an outline of the thesis. The theory presented in this chapter is based on Thompson \etal\ \cite{thompson2008interferometry} and some online sources, notably presentations given in a workshop \cite{Perley2010} organised by the \textit{National Radio Astronomy Observatory}(NRAO), and a course \cite{Condon} given by the same organisation.	\subsection{\textit{Compression}} In the first batch of experiments (refer to Figure \ref{fig:res-batch1}), switching off \textit{channel interleaving} while compression is disabled (experiment 1.2) reduces the gridder performance for the large convolution functions (31$\times$31 and wider), by around 10G. No change in performance resulted for smaller convolution functions. On enabling compression, the total rate for interleaved channels goes to a maximum of 280G and for non-interleaved channels goes to 175G. This implies a 3-fold performance increase for \textit{interleaved channels} and 2-fold increase for \textit{non-interleaved channels}. The real gridding rate goes down as compression is enabled, but as argued in section \ref{sec:realrate} it goes to a steady value. From these results, it can be deduced that performance of Romein's algorithm without compression has some dependence on the records that would be compressed if compression is enabled. The methods implemented in this thesis tackle such "compressible" records in a way to obtain higher performance. Unfortunately, the claimed total rates are not likely to be achieved in realistic scenarios where proper \textit{w-projection} is used. In section \ref{sec:compression} it was discussed that the probability of compression is not evenly distributed over the UV-grid but is more prominent for short baselines. Short baselines tend to have small convolution functions, and thus the highest probability for compression is for records with the least computational exigencies. This argument is well supported by the reported results of experiment batch 5. Figures \ref{fig:res-compressionratio} and \ref{fig:res-avg-support} show that for all the runs made using proper \textit{w-projection}, the average convolution function size for compressed records is smaller than that for really gridded records. Consequently, the \textit{compression} ratio in terms of grid point updates\footnote{This controls the total gridding rate.} is smaller than the respective record \textit{compression} ratio. This reduction reached a factor of 10 for some runs. This implies that performance gains delivered by \textit{compression} can be severely degraded. This does not necessarily make \textit{compression} ineffective since as shown in Figure \ref{fig:res-overall-gridderrate}, Data Set 2 is gridded using Configuration 1 at a total rate of 107G. One notes that when there are no records to compress there is no performance loss against a scenario with \textit{compression} disabled. \textit{Compression} can either not affect performance or enhance it. Integration time effects \textit{compression}. Data Set 2 has an integration time roughly 3$\times$ shorter than Data Set 1, and results depicted in Figure \ref{fig:res-compressionratio} reveal the highly different \textit{compression} ratios between Data Set 1 and Data Set 2. This is in-line with the arguments given in section \ref{sec:compression}. A final note on \textit{compression} regards \textit{channel interleaving}. From the results reported on the first batch of experiments, it is clear that \textit{compression} occurred over channels. This shows that \textit{compression} can also deliver performance in multi-frequency image synthesis. In such synthesis, channels are gridded over the same grid. The developed imaging tool does not support such a feature, but it is a good point to note for the future. The good news about \textit{compression} over channels is that the probability of occurrence should be distributed evenly over the grid. It is likely to be dependent on channel frequency bandwidth, whereby the narrower the channel frequency bandwidth, the higher is the probability for \textit{compression}. Worth recalling that, for a wide field of view, channel bandwidth has to be as narrow as possible (refer to section \ref{sec:channelbadwidth}). \subsection{Real gridding rate and gridder scalability} \label{sec:realrate} It is deduced from the second and third batch of results (Figures \ref{fig:res-deltau-set0}- \ref{fig:res-real}) that there is little effect on the gridder real performance when varying the oversampling factor or the UV-grid sampling interval. Figure \ref{fig:res-real} gives the average rate obtained from these experiments per convolution function size and reports the total variations in performance. The biggest fluctuations happened for the $15\times15$ convolution function whereby a maximum rate variation of 16G resulted. For all the others, the maximum variation was less than 10G. It is also deduced that the gridder is scalable in terms of convolution function size. The gridder, grids convolution functions of size larger or equal to $31\times31$ at a nearly constant rate of around 50G. There is clearly a degradation in performance for convolution function of size less than $31\times31$. However, this does not constitute a real issue since they are relatively light in terms of computation. Rates are nearly maintained when proper \textit{w-projection} is enabled. Figure \ref{fig:res-overall-gridderrate} reports that, for the heavy-weight Configuration 2, the gridder rate was around 49G for the two data sets. It is slightly less than the gridding rate achieved by convolution functions larger or equal to $31\times31$, but much higher than the gridding rate achieved for some of the small convolution functions ($7\times7$ and $23\times23$). In view that proper \textit{w-projection} uses many convolution function sizes and that the local chunk length (refer to section \ref{sec:localchunks}) is variable, then, this rate can be considered as acceptable. Performance achieved using the light-weight Configuration 1 is also acceptable though much less than Configuration 2. Given that, in this configuration, all convolution sizes are less than $31\times31$ pixels, then such a drop is expected since these sizes tend to be gridded at a lower rate. \subsection{Number of polarisations} \label{sec:res-polarizations} Results of experiment batch 4 reported in Figures \ref{fig:res-polarizations} and \ref{fig:res-pol-comp}, show that the gridder does not down-scale with the number of polarisations being gridded. Gridding dual or single polarised records require half or quarter of grid point updates respectively than gridding quad-polarised records. If the gridding rate (in grid point updates/sec) remains constant over the number of polarisations, then, records are expected as to be gridded at a double or quadruple rate respectively when compared to quad-polarisation gridding. Instead, as reported in Figure \ref{fig:res-pol-comp} the record gridding rate increased only by a factor of 1.1$\times$ for single-polarisation gridding of large convolution functions. The culprit is believed to be the retrieval of convolution function data that is discussed in section \ref{sec:lim}. \subsection{Main limiting factor of the gridder} \label{sec:lim} It is claimed that the main limiting factor of the gridder is the retrieval of the convolution function numerical data from texture. This happens for each grid point update. Romein \cite{Romein2012}, argues that the main limiting factor of his scenario is atomic operations. This behaviour is not observer in the implementation presented in this thesis, since there is little change in performance when the oversampling factor is varied. A reduction in oversampling rate results in a reduction in the number of records that are gridded without causing any atomic commits, since more records get compressed. Thus for each grid point update the likelihood of an atomic commit is increased. If atomic operations were the main limiting factor, the real gridding rate should drastically decrease with decreasing oversampling factor. Instead, the rate increases for large convolution functions (refer to Figures \ref{fig:res-sam-set0}, \ref{fig:res-sam-set1}, \ref{fig:res-sam-set2}). For smaller convolution functions, the rate does decrease but not in the order that would be expected. A similar argument can be made for the polarisation results (batch 4). Decreasing polarisations reduces atomic operations per gridded record in proportion to the decrease in the number of polarisations. No substantial increase in record gridding rate resulted when the number of polarisations is decreased. The polarisation results also rule out the floating point operations required for grid point update as a main limiting factor. The payload is constant to eight flops per grid point and independent on the number of polarisations. In view of the heavy reduction in the gridding rate expressed in grid point updates per second, it cannot be the main limiting factor. The remaining possible culprits are the retrieval of the convolution function from texture, the gridding logic and access to shared memory. When \textit{compression} is disabled (see Figure \ref{fig:res-batch1}), the real rate increases drastically (around 30G). Payload of shared memory access and gridding logic is fixed for enabled or disabled \textit{compression}. The only variant is the texture performance since it is a cache. When \textit{compression} is disabled the likelihood of a cache hit is increased. This is because the probability that the same numerical data is used by subsequent records is larger than zero. Hence, the claim that the main limiting factor is the retrieval of convolution function data from the texture is proved. \subsection{\textit{WImager} preparation phase performance} Figure \ref{fig:res-overall-rate} shows that the preparation phase prepares records at a fast rate. Its impact on the overall performance of the \textit{WImager} algorithm is dependent on the computational intensity required by the gridder. The higher the computational intensity of the gridder phase, the lower is the overall impact of the preparation phase. Figure \ref{fig:res-kernel-time} visualises the point. For the heavy-weight Configuration 2, the preparation phase execution time is negligible when compared to the execution time of the gridder. On the other hand, imaging over Configuration 1 resulted in a preparation phase execution time larger than the gridding time. One must bear in mind that the preparation phase's main job is to reduce logic from the gridding phase. If the logic is integrated in the gridding phase, the total execution time of the \textit{WImager} algorithm would increase. \subsection{Overall performance} Overall performance of the \textit{mt-imager} is reported in Figure \ref{fig:timeline}. Results for the multi-GPU scenario are ignored in this section but discussed in section \ref{sec:multiGPU}. Figure \ref{fig:timeline} reveals that \textit{mt-imager} synthesised Data Set 2 over Configuration 2 (the hardest of all runs) 94$\times$ faster than \textit{lwimager}. \textbf{This result shows that the main thesis objective of developing a high-performance imaging tool has been achieved.} This gain is not sustained for all runs. It is a side effect of high performance. Computation is so efficient that the loading of data from disk\footnote{As per section \ref{sec:visibilitymanager}, data is loaded through the casacore ms API which can affect the data loading time.} is a significant limiting factor. During this time, the GPU has to wait for the first chunk of data. Worst case occurred for the simplest run (Data Set 1 over Configuration 1), where only a 6.2$\times$ gain was obtained. Loading of data in this run took most of the time\footnote{Exit time is ignored in this discussion.}. It should be stated that for the simplest run (Data Set 1 over Configuration 1), the generation of the 200 convolution functions over the GPU finished nearly at the same time when loading of data was ready. This implies that generation of convolution functions might sometimes limit the performance further. One notes the remarkably small time interval, shown in red, for all runs. During this time, the \textit{Visibility Manager} makes the last preparations for the first chunk of data after that all data has been loaded from disk. This time interval is short because the \textit{Visibility Manager} did most of the preparation work while the system is still loading data. It does not make miracles, and for the large Data Set 1 it requires a substantial amount of extra time to prepare the other data chunks. More in-depth analyses reveal that this extra time is needed to sort and convert visibility data (that is the sequence defined as $\{\vis_i\}_p$ is Table \ref{tab:legend}). Visibility data has the highest memory consumption. This extra time can also impose limits on the performance of the imaging tool. The \textit{Visibility Manager} might not supply data chunks at rates faster than the processing of the chunks over the GPU. A case in point is Data Set 2 imaged using Configuration 1, where most of the GPU work is done while the \textit{Visibility Manager} is preparing data. Results show that the processing time of the \textit{Visibility Manager} is independent of configuration but mostly dependent on the data set being imaged. This is an expected result. A final observation is the excessive time the imaging tool takes to exit. It is marked in Figure \ref{fig:timeline} as \textit{exit time}. This time interval is a waste of time since by then, all images are finalised and saved to disk. For the multi-GPU scenario, it amounted to 10 seconds! Detailed analysis revealed that most of this time is consumed by the CUDA Runtime API to reset GPU devices. It is yet unclear who is the culprit, whether a limitation of the CUDA Runtime API, GPUs, or something else. \subsection{Multi-GPU scalability} \label{sec:multiGPU} When imaging over two GPUs (refer to Figure \ref{fig:timeline}), only a performance gain of $1.2\times$ against the GTX670\footnote{Comparison is made against the GTX670 GPU because it is less powerful than the GTX680.} run was obtained. The rather low value is the result of the exceptionally strong performance already obtained by imaging over 1 GPU. Performance gains obtained from imaging over more than 1 GPU are limited by the time consumed to load data from disk. If the time to load data and the \textit{exit time} are ignored, a speed-up of $1.8\times$ results. Thus, \textit{mt-imager} is scalable over GPUs. Nevertheless, the loading of data from disk, limits the gains severely. \subsection{Summary} Table \ref{tab:maintopics} summarises the main performance topics reported in this chapter. \begin{table}[H] \centering \begin{tabularx}{\textwidth}{LL} \hline \textit{Topic} & \textit{Comment} \\ \hline Overall performance & Nearly $100\times$ faster then \textit{lwimager}. \\ Main limiting factor & Loading of data from disk. \\ Other limiting factor & \textit{mt-imager} takes substantial time to exit. \\ Scalability over GPUs & \textit{mt-imager} is scalable over GPU, but most gains are hindered by the main limiting factor. \\ \textit{WImager} gridder performance & Real gridding rate of ~50 Giga grid point updates/sec for most computationally intensive scenario. \\ \textit{Compression} performance & Performance obtained from \textit{compression} varies depending on imaging configuration and data set. 3-fold increase in gridding performance were obtained in particular runs.\\ Gridder main limiting factor & Retrieval of convolution function numeric data from GPU memory.\\ \hline \end{tabularx} \caption{\textit{mt-imager} performance summary} \label{tab:maintopics} \end{table} \chapter{Conclusion} \label{chap:conclusion} In this thesis, a new high-performance imaging synthesis tool for radio interferometry was developed. The tool which is called \textit{malta-imager} or \textit{mt-imager} exploits the computational power delivered by GPUs, to achieve unprecedented high performance. The backbone handling numerical calculations was generalised and a new framework was developed called the \textit{General Array Framework} (GAFW). Test cases presented in this thesis show that the imaging tool is able to synthesis images nearly $100\times$ faster than a common CPU based imaging tool. This clearly shows that the thesis main objective, which is the development of a high-performance imaging tool, has been achieved in full.	14	3	1403.4209	A radio interferometer indirectly measures the intensity distribution of the sky over the celestial sphere. Since measurements are made over an irregularly sampled Fourier plane, synthesising an intensity image from interferometric measurements requires substantial processing. Furthermore there are distortions that have to be corrected. In this thesis, a new high-performance image synthesis tool (imaging tool) for radio interferometry is developed. Implemented in C++ and CUDA, the imaging tool achieves unprecedented performance by means of Graphics Processing Units (GPUs). The imaging tool is divided into several components, and the back-end handling numerical calculations is generalised in a new framework. A new feature termed compression arbitrarily increases the performance of an already highly efficient GPU-based implementation of the w-projection algorithm. Compression takes advantage of the behaviour of oversampled convolution functions and the baseline trajectories. A CPU-based component prepares data for the GPU which is multi-threaded to ensure maximum use of modern multi-core CPUs. Best performance can only be achieved if all hardware components in a system do work in parallel. The imaging tool is designed such that disk I/O and work on CPU and GPUs is done concurrently. Test cases show that the imaging tool performs nearly 100$\times$ faster than another general CPU-based imaging tool. Unfortunately, the tool is limited in use since deconvolution and A-projection are not yet supported. It is also limited by GPU memory. Future work will implement deconvolution and A-projection, whilst finding ways of overcoming the memory limitation.	false	[ "modern multi-core CPUs", "imaging tool", "GPU memory", "maximum use", "substantial processing", "GPUs", "several components", "unprecedented performance", "Graphics Processing Units", "CPU", "numerical calculations", "oversampled convolution functions", "Best performance", "GPU", "interferometric measurements", "parallel", "use", "another general CPU-based imaging tool", "a new high-performance image synthesis tool", "Future work" ]	10.954505	3.904649	171
12371635	[ "Hoffmann, K.", "Bel, J.", "Gaztañaga, E.", "Crocce, M.", "Fosalba, P.", "Castander, F. J." ]	2015MNRAS.447.1724H	[ "Measuring the growth of matter fluctuations with third-order galaxy correlations" ]	63	[ "Institut de Ciències de l'Espai (ICE, IEEC/CSIC), E-08193, Barcelona, Spain", "INAF - Osservatorio Astronomico di Brera, Via Brera 28, 20122 Milano, via E. Bianchi 46, I-23807 Merate, Italy; Aix Marseille Université, CNRS, Centre de Physique Théorique, UMR 7332, F-13288 Marseille, France; Université de Toulon, CNRS, CPT, UMR 7332, F-83957 La Garde, France", "Institut de Ciències de l'Espai (ICE, IEEC/CSIC), E-08193, Barcelona, Spain", "Institut de Ciències de l'Espai (ICE, IEEC/CSIC), E-08193, Barcelona, Spain", "Institut de Ciències de l'Espai (ICE, IEEC/CSIC), E-08193, Barcelona, Spain", "Institut de Ciències de l'Espai (ICE, IEEC/CSIC), E-08193, Barcelona, Spain" ]	[ "2014MNRAS.443.2874M", "2015MNRAS.447..646C", "2015MNRAS.450.1674H", "2015MNRAS.450.1836K", "2015MNRAS.451.4029K", "2015MNRAS.453..259B", "2015MNRAS.453.1513C", "2015PhRvD..91d3530S", "2016MNRAS.462...35P", "2016arXiv160603892V", "2017A&A...599A..79P", "2017A&A...604A.133M", "2017ApJ...836...54M", "2017MNRAS.465.2225H", "2017MNRAS.466.1444C", "2018JCAP...09..008L", "2018MNRAS.473.3051I", "2018MNRAS.476..814H", "2018MNRAS.481.5189B", "2018PhR...733....1D", "2019A&A...624A..30J", "2019ApJ...871L..13F", "2019ApJ...882..166M", "2019MNRAS.482.3341C", "2019MNRAS.488.5452C", "2019arXiv191206906L", "2020A&A...637A.100W", "2020A&A...642A..83D", "2020A&A...642A.200V", "2020A&C....3200391T", "2020ApJ...890...78A", "2020ApJ...905..127D", "2020JCAP...07..043K", "2020JCAP...08..007D", "2021A&A...645A.104A", "2021A&A...645A.105G", "2021A&A...646A.129J", "2021A&A...647A.124H", "2021A&A...647A.185G", "2021A&A...649A..99S", "2021A&A...650A.113B", "2021A&A...650A.148S", "2021A&A...654A..76F", "2021ApJ...919...69C", "2021ApJ...919..144M", "2021MNRAS.503.4964G", "2021MNRAS.504.1452V", "2021MNRAS.506.4344N", "2022A&A...661A..70L", "2022A&A...664A.170V", "2022ApJ...940..115C", "2022MNRAS.515.2305S", "2022MNRAS.517.4827G", "2022PhRvD.105b3515S", "2023A&A...675A.189D", "2023AJ....166...22G", "2023MNRAS.523.3133S", "2024ApJ...964..191B", "2024JCAP...02..024D", "2024arXiv240104687A", "2024arXiv240507904P", "2024arXiv240514818H", "2024arXiv240605122L" ]	[ "astronomy" ]	2	[ "cosmology: miscellaneous", "cosmology: observations", "cosmology: theory", "dark energy", "dark matter", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1977ApJ...217..385G", "1985ApJ...292..371D", "1986ApJ...311....6G", "1987MNRAS.227....1K", "1992ApJ...392....1B", "1992ApJ...398L..17G", "1993ApJ...413..447F", "1994ApJ...425..392F", "1994ApJ...433....1B", "1994ApJ...437L..13G", "1994MNRAS.268..913G", "1994PhRvL..73..215F", "1996A&A...312...11B", "1997MNRAS.284..189M", "1998AJ....116.1009R", "1998ApJ...508..483W", "1998MNRAS.299.1097S", "1998MNRAS.300L..35S", "1999ApJ...514....7K", "1999ApJ...517..565P", "1999ApJ...521L..83F", "2000ApJ...530...36B", "2001ApJ...548...47G", "2001MNRAS.323....1S", "2001PhRvL..86.1434F", "2002MNRAS.331...13G", "2002MNRAS.333..443B", "2002MNRAS.335..432V", "2002PhR...367....1B", "2004ApJ...614..527Z", "2004PhRvD..69d4005L", "2005A&A...442..801M", "2005MNRAS.361..824G", "2005MNRAS.362.1363P", "2005MNRAS.364..620G", "2005MNRAS.364.1105S", "2006MNRAS.365..214S", "2006PhRvD..74b3522S", "2007MNRAS.381..573R", "2008A&A...487....7M", "2008MNRAS.387..921A", "2008Natur.451..541G", "2009JCAP...10..004S", "2009MNRAS.393.1183C", "2009MNRAS.396...19N", "2011ApJ...731..102K", "2011ApJ...737...97M", "2011ApJ...739...85M", "2011MNRAS.415..383M", "2011MNRAS.415.2876B", "2012ApJ...750...37J", "2012MNRAS.420.2102S", "2012MNRAS.420.3469P", "2012MNRAS.424..971B", "2012MNRAS.424.2339T", "2012MNRAS.426.2719R", "2012PhRvD..85h3509C", "2012PhRvD..86h3540B", "2013A&A...557A..54D", "2013MNRAS.430..924C", "2013MNRAS.430.2476S", "2013MNRAS.432.2654M", "2014A&A...571A..16P", "2014JCAP...05..042S", "2014MNRAS.443.2874M", "2015MNRAS.447..646C", "2015MNRAS.447.1319F", "2015MNRAS.448.2987F", "2015MNRAS.453.1513C" ]	[ "10.1093/mnras/stu2492", "10.48550/arXiv.1403.1259" ]	1403	1403.1259.txt	Evidence that the expansion of the Universe is accelerating \citep{de1,de2} has revived the cosmological constant $\Lambda$, originally introduced by Einstein as an unknown fluid which may engine the observed dynamics of the Universe. Alternative explanations for the accelerated expansion could involve a modification of the gravitational laws on cosmological scales. Since these modifications of gravity can mimic well the observed accelerated expansion it is difficult to just rely on the cosmological background (i.e. the overall dynamics of the Universe) in order to verify which model is correct. However, alternative gravitational laws change the way matter fluctuations grow during the expansion history of our Universe. Measuring the growth of matter fluctuations could therefore be a powerful tool to distinguish between cosmological models \citep[see e.g.][]{GL, Lue04, rossetal07, S&P09, C&G09, SP&R12, reidetal12, contrerasetal13, guzzoetal08, delatorreetal13, SB&M}. On this basis, the goal of several future and ongoing cosmological surveys, such as BOSS\footnote{https://www.sdss3.org/surveys/boss.php}, DES\footnote{www.darkenergysurvey.org}, MS-DESI\footnote{desi.lbl.gov}, PAU\footnote{www.pausurvey.org}, VIPERS\footnote{http://vipers.inaf.it} or Euclid\footnote{www.euclid-ec.org}, is to measure the growth of matter fluctuations. This can be achieved by combining several observables, such as weak gravitational lensing, cluster abundance or redshift space distortions. Higher-order correlations in the galaxy distribution provide additional observables which also allow for proving the growth equation beyond linear theory from observations \citep[e.g. see][]{bernardeau02}. Furthermore, higher-order correlations can be used to test the nature of the initial conditions and improve the signal-to-noise in recovering cosmological parameters \citep[e.g.][]{Sefusatti06}. The relative simplicity of the fundamental predictions about amplitude and scaling of clustering statistics, must not make us overlook the fundamental difficulty that hampers large scale structure studies. The perfect, continuous (dark matter) fluid in terms of which we model the large-scale distribution of matter cannot be directly observed. Let's imagine that we are able to locate in the Universe all existing galaxies and that we know with an infinite precision their masses. Without any knowledge of how luminous galaxies trace the underlying continuous distribution of matter, even this ultimate galaxy sample would be of limited use. The problem of unveiling how the density fields of galaxies and mass map into each other is the so called galaxy biasing. Knowledge of galaxy bias, and therefore galaxy formation, can greatly improve our cosmological inferences from observations. A common approach to model galaxy bias consists in describing the mapping between the fields of mass and galaxy density fluctuations ($\delta_{dm}$ and $\delta_g$ respectively) by a deterministic local function $F$. This function can be approximated by its Taylor expansion if we smooth the density field on scales that are sufficiently large to ensure that fluctuations are small, \begin{equation} \delta_g=F[\delta_{dm}] \simeq \sum_{i=0}^{N}\frac{b_i}{i!}\delta_{dm}^i, \label{eq:biasfunction} \end{equation} % where $b_i$ are the bias coefficients. It has been shown that, in this large scale limit, such a local transformation preserves the hierarchical properties of matter statistics \citep{FG}. There is now convincing evidences about the non-linear character of the bias function \citep{ga92,mar05, gnbc, mar08, Kovac}. Since we only want to study correlations up to third order, in this paper we shall consider bias coefficients up to second order, i.e. $b_1$ and $b_2$, which is expected to be sufficient at the leading order \citep{FG}. However, one of the goals of this paper is to investigate at which scale and halo mass range this expectation is fulfilled. To study the statistical properties of the matter field we need to find the most likely value for the coefficients $b_i$. A general approach aims at extracting them from redshift surveys using higher-order statistics. If the initial perturbations are Gaussian and if the shape of third-order statistics are correctly described by results of the weakly non-linear perturbation theory, then one can fix the amplitude of $b_i$ up to second order in a way which is independent from the overall amplitude of clustering (e.g. $\sigma_8$) and depends only on the shape of the linear power spectrum. This has been shown by several authors using the the skewness $S_3$ \citep{gazta94, gaztafrie94}, the bispectrum \citep{Fry94,gaztafrie94,scoc98,feld01,Verde}, the three-point correlation function $Q$, \citep{gnbc, gaztascoc, panszapudi,marin11,mcbride11,marin13}, and the two-point cumulants $C_{12}$ \citep{bernardeau96,Szapudi98,GFC,bm}. Recently, \citet{bm} demonstrated that it is possible to use these higher-order correlations to constrain bias and fundamental properties of the underlying matter field using a combination of $S_3$ and $C_{12}$, which they call $\tau =3C_{12}- 2S_3$. The main goal of this paper is to present for the first time a comparison of the bias derived from this new $\tau$ method with that of $Q$, using the same simulations and halo samples. We also show that, with a new approach, the growth of matter fluctuations can be measured directly from observations by getting rid of galaxy bias and without requiring any modelling of the underlying matter distribution. Despite the fact that in the present analysis we only consider real-space observables (not affected by redshift-space distortions) we argue that, as long as reduced third-order statistics are only weakly affected by redshift-space distortions (for a broad range of masses, see Fig. \ref{fig:q3pz}), the proposed method appears to be applicable on redshift galaxy surveys. This analysis is based on the new MICE-GC simulation and extends its validation presented recently by \citet{mice1, mice2, mice3}. In Section \ref{sec:micegc} we present the simulation on which our work relies. Our estimators for both, the bias and the growth of matter fluctuations, are introduced in Section \ref{sec:methods}. We present our results in Section \ref{sec:results} and a summary of the work can be found in Section \ref{sec:disc} together with our conclusions. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% SIMULATION (Section II) %% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%	\label{sec:disc} The amplitude of the transverse (or projected) two-point correlation of matter density fluctuations allows us to measure the growth factor $D$, which can be used as a verification tool for cosmological models. Galaxies (in our study represented by haloes) are biased tracers of the full matter field as their two-point correlation at large scales is shifted by a constant bias factor $b$ with respect to the matter two-point correlation. This bias factor is fully degenerate with $D$. The reduced matter and galaxy third-order statistics are independent of $D$, while the galaxy versions are sensitive to $b$. Combining second- and third-order statistics could therefore enable us to break the growth-bias degeneracy, if the difference between the effective linear bias $b_1$ probed by both statistics is smaller than the errors required for the growth measurements. In this paper we have tested these assumptions and verified how well we can recover the true growth of the new MICE-GC $\Lambda$CDM simulation \citep{mice1, mice2, mice3} with them. We also further validate the MICE-GC simulation by comparing the linear growth with the two-point matter correlation (Fig.2 and 3) and the different third-order statistics of the matter field to non-linear perturbation theory predictions (Fig.5 and 8). In particular, previous analysis \citep{GFC} found a mismatch between simulations and predictions for $C_{12}$ \citep{bernardeau96}, which we find here to originate from neglecting one of the smoothing terms (i.e. $\beta_R$ in Eq.28). After taking this into account, the MICE-GC simulations agrees well with predictions for all redshifts (see Fig.8). This, therefore, provides a validation of the approach adopted by \citet{bm} to measure the linear galaxy bias using only galaxy clustering. The main goal of this paper is to compare bias (and the resulting growth) measurements from two different third-order statisticics proposed in the literature. One uses the reduced three-point correlation $Q$ while the other uses a combination of the skewness $S_3$ and two-point third-order correlator, $C_{12}$, which is called $\tau\equiv 3C_{12}-2S_3$. We estimated these quantities from density fields of matter in the MICE-GC simulation and those of haloes in different mass samples, expanding previous studies significantly to a wider range of masses (between $5.8\times 10^{12}$ and $5\times 10^{14}$ $h^{-1}$M$_\odot$) and redshifts (between $0$ and $1.2$) with values of the linear bias $b_1$ between $0.9$ and $4$. Our results in Fig. \ref{fig:results4} show that the linear bias from $Q$, $b_Q$, systematically over estimates the linear bias from the two-point correlation, $b_{\xi}$, by roughly $20-30\%$ at all mass and redshift ranges, whereas the linear bias from $\tau$, $b_\tau$, seems to be an unbiased estimator at the price of decreased precision. Non-local contributions to galaxy bias, like tidal effects, are anisotropic and therefore could be more important for $b_Q$ than for $b_\tau$, as $\tau$ is isotropic (i.e. it comes from higher-order one- and two-point correlations, while $Q$ comes from three points). In Fig. \ref{fig:q3pz} we illustrate the different impacts of the local bias model and the non-local model of \cite{chan12} with $\gamma_2=2(b_1-1.43)/7$ on $Q$ (dashed and solid lines respectively). The non-local model seems to approximate $Q$ measurements from halo samples better but there are still some discrepancies that we will explore in a separate analysis. Besides non-local contributions to the bias model, further reasons for the difference between $b_Q$ and $b_{\tau}$ might be that non-linear terms in the bias function and the matter field have different impacts on $Q$ and $\tau$. In addition we found that estimations of the quadratic bias parameter $c_2$ from $Q$ and $\tau$ can also differ significantly from each other. Understanding the differences between $b_{\xi}$, $b_Q$ and $b_{\tau}$ is crucial for constraining cosmological models with observed third-order halo statistics. We will therefore deepen our analysis in a second paper by studying bias from halo-matter-matter statistics, direct analysis of the halo versus matter fluctuations and predictions from the peak-background split model to disentangle between non-linear and non-local effects on the different estimators. For measuring the growth factor $D$ we have introduced a new method. This new method uses the bias ratio $\hat{b}(z)=b(z)/b(z_0)$, derived directly from halo density fluctuations with reduced third-order statistics. Its main advantage with respect to the approach of measuring $b(z)$ and $b(z_0)$ separately is that it does not require the modelling of (third-order) dark matter statistics. Instead, it works with the hypothesis that \begin{enumerate} \item the reduced dark matter three-point statistics is independent of redshift $z$ \item the bias ratio $\hat{b}(z)=b(z)/b(z_0)$ from two- and three-point statistics is equal. \end{enumerate} The first assumption was tested in this study numerically, while the validity of the second follows directly from our bias comparison. In general the comparison between $D$ from perturbation theory with measurements from our new method and the standard approach reveals a good agreement. In the case of $Q$ we explain this result by a cancellation of the multiplicative factor by which $b_Q$ is shifted away from $b_{\xi}$ in the bias ratio $\hat b_Q$. The growth factor measured with $\tau$ has larger errors than the results from $Q$ as a consequence of the larger errors in the bias estimation. \begin{figure} \centering \includegraphics[width=85mm]{./plots/q3dm24RSDzr.eps} \caption{$Q$ for dark matter (dotted) and for halo samples (symbols) with two different mass thresholds: $b_1 =b_\xi\simeq 1.09$ (blue) and $b_1 =b_\xi \simeq 1.83$ (red). We compare results in real space (filled triangles) and redshift space (open circles), which agree within the errors on these large scales ($r_{12} = r_{13}/2 = 24$ $h^{-1}$Mpc at z=0). Predictions are shown for both: the local bias model (dashed lines) and non-local bias model (continuous). In both cases we have fixed $b_1=b_\xi$ and fit for $c_2$. } \label{fig:q3pz} \end{figure} Our analysis shows that the new way to measure the growth factor from bias ratios is competitive with the method based on two separate bias measurements. While having larger errors the new method has the advantage of requiring much weaker assumptions on dark matter correlations than the standard method and therefore provides an almost model independent way to probe the growth factor of dark matter fluctuations in the Universe. We demonstrated that besides the growth factor, $D$, the growth rate of matter, $f$, can also be directly measured from the galaxy (or halo) density fields with bias ratios from third-order statistics. This provides an alternative method to derive the growth rate, which is usually obtained from velocity distortions probed by the anisotropy of the two-point correlation function (RSD). The typical errors found on SDSS, BOSS and WiggleZ using RSD are around 15-20\% \citep{C&G09,Blake2011,Tojeiro2012}, which are comparable to the ones we find in Fig. \ref{fig:results10} when considering the high redshift bins (20\%). Given that the two methods explored here use different information from higher-orders correlation ($Q$ uses the shape, while $\tau$ uses collapse configurations) one can reasonably guess that the two methods are not strongly correlated. So a possible strategy would be to use the $Q$ method (more precise) to measure the (velocity) growth rate and, in parallel, to use the $\tau$ method to extract the growth factor. This would help to break degeneracies between cosmological parameters in different gravitational frameworks. Our analysis is performed in real space to have clean conditions for comparing different bias and growth estimates. This is a good approximation for the reduced higher-order correlations on the large scales considered in this study, as measurements in redshifts space always seem to be within one sigma error of the corresponding real space result (see Fig. \ref{fig:q3pz}). Note how the small, but systematic, distortions in redshifts space seem to agree even better with the local bias model than in real space on the largest scales. Applying the methods described above to obtain accurate bias and growth measurements from observations will require additional treatment of redshifts space distortions or projection effects. Two possible paths could be followed. In a three dimensional analysis redshifts space distortions need to be modeled \citep[e.g.][]{gaztascoc}. The projected three-point correlation can also be studied separated by in redshift bins \citep{friga1999,Buchalter,zheng}. Both ways will result in larger errors, but we do not expect this to be a limitation because our error budget is totally dominated by the uncertainty in the bias. A more detailed study of this issue is beyond the scope of this paper and will be presented elsewhere. Mock observations, like the galaxy MICE catalogues \citep[see][]{mice2,Carretero2014} should be used to test the validity of such growth measurements under more realistic conditions. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% ACKNOWLEDGEMENT %% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%	14	3	1403.1259	Measurements of the linear growth factor D at different redshifts z are key to distinguish among cosmological models. One can estimate the derivative dD(z)/dln (1 + z) from redshift space measurements of the 3D anisotropic galaxy two-point correlation ξ(z), but the degeneracy of its transverse (or projected) component with galaxy bias b, i.e. ξ<SUB>⊥</SUB>(z) ∝ D<SUP>2</SUP>(z)b<SUP>2</SUP>(z), introduces large errors in the growth measurement. Here, we present a comparison between two methods which breaks this degeneracy by combining second- and third-order statistics. One uses the shape of the reduced three-point correlation and the other a combination of third-order one- and two-point cumulants. These methods use the fact that, for Gaussian initial conditions and scales larger than 20 h<SUP>-1</SUP> Mpc, the reduced third-order matter correlations are independent of redshift (and therefore of the growth factor), while the third-order galaxy correlations depend on b. We use matter and halo catalogues from the MICE-GC simulation to test how well we can recover b(z) and therefore D(z) with these methods in 3D real space. We also present a new approach, which enables us to measure D directly from the redshift evolution of the second- and third-order galaxy correlations without the need of modelling matter correlations. For haloes with masses lower than 10<SUP>14</SUP> h<SUP>-1</SUP> M<SUB>⊙</SUB>, we find 10 per cent deviations between the different estimates of D, which are comparable to current observational errors. At higher masses, we find larger differences that can probably be attributed to the breakdown of the bias model and non-Poissonian shot noise.	false	[ "matter correlations", "non-Poissonian shot noise", "redshift space measurements", "different redshifts", "third", "galaxy bias b", "3D real space", "non-Poissonian", "current observational errors", "large errors", "cosmological models", "Measurements", "redshift", "the reduced third-order matter correlations", "the third-order galaxy correlations", "matter", "D", "larger differences", "the second- and third-order galaxy correlations", "second- and third-order statistics" ]	11.891466	2.537554	161
437639	[ "Cox, N. L. J.", "Patat, F." ]	2014A&A...565A..61C	[ "Dense molecular clouds in the SN 2008fp host galaxy" ]	21	[ "Instituut voor Sterrenkunde, KU Leuven, Celestijnenlaan 200 D, 3001, Leuven, Belgium", "European Organization for Astronomical Research in the Southern Hemisphere (ESO), Karl Schwarzschild-Str. 2, 85748, Garching bei München, Germany" ]	[ "2014ApJ...792..106W", "2014ApJ...795L...4K", "2015A&A...577A..53P", "2015ApJ...799..197R", "2015JChPh.143g4302R", "2015MNRAS.451.4104J", "2015MmSAI..86..541Z", "2015PASP..127..223R", "2016ApJ...828...24P", "2017A&A...606A..76C", "2017A&A...606A.109M", "2017ApJ...835..100M", "2017ApJ...836...13H", "2017ApJ...836...88Z", "2017suex.book.....B", "2018EPJD...72..134V", "2018RSPTA.37670145V", "2019ApJ...883..122L", "2020ApJ...888...93G", "2022ApJ...939...18Y", "2024A&A...681A...6F" ]	[ "astronomy" ]	0	[ "supernovae: individual: SN 2008fp", "ISM: lines and bands", "ISM: molecules", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1972ARA&A..10..305T", "1975ApJ...196..261S", "1980ARA&A..18..489W", "1982ApJ...257..125F", "1983ApJ...272..509A", "1985ApJ...297..119M", "1985ApJ...299..852D", "1987AJ.....94..651R", "1987MNRAS.224..299S", "1987MNRAS.227P...1H", "1988PASA....7..527P", "1989A&A...215...21D", "1989MNRAS.241..575C", "1990AJ.....99.1476S", "1991A&A...251..625G", "1991ApJ...381L..17C", "1991ApJ...383L..41M", "1991rc3..book.....D", "1992ApJ...386..562W", "1992dge..book.....W", "1994ApJ...420L..71B", "1994ApJ...424..748C", "1998ApJ...500..525S", "1999A&A...351..657G", "1999ApJ...512..511M", "2001ApJ...553..267M", "2003AJ....126.2268W", "2003ApJ...582..823O", "2003ApJ...584..339T", "2003ApJ...595..235A", "2004ApJ...614..658J", "2005A&A...429..559S", "2005ApJ...633..986P", "2006A&A...447..991C", "2006ARA&A..44..367S", "2006ApJ...647L..29Y", "2006ApJ...649..788R", "2006ApJS..165..138W", "2006MNRAS.367.1478H", "2007ApJS..168...58S", "2007Sci...317..924P", "2008A&A...480..133L", "2008A&A...480L..13C", "2008A&A...485L...9C", "2008A&A...492L...5C", "2008AJ....136..994L", "2008ApJ...687.1075S", "2008CBET.1506....1P", "2008CBET.1509....1W", "2008MNRAS.383L..30E", "2008MNRAS.390.1733S", "2009A&A...508..229P", "2009ApJ...693..207B", "2009ApJ...700.1299J", "2009ApJ...702.1157S", "2009ApJ...707..916F", "2010A&A...514A..78P", "2010AJ....140.2036S", "2010MNRAS.404.1321W", "2011A&A...533A.129V", "2011AJ....142..156S", "2011ApJ...727...33F", "2011ApJ...728...36R", "2011Sci...333..856S", "2012ApJ...748L..11R", "2012MNRAS.421.1325L", "2013A&A...549A..62P", "2013ApJ...779...38P", "2013MNRAS.428.1107W", "2013MNRAS.428.2198S", "2014IAUS..297..106C" ]	[ "10.1051/0004-6361/201219143", "10.48550/arXiv.1403.4386" ]	1403	1403.4386_arXiv.txt	\label{sec:intro} Multi-epoch high-resolution optical spectroscopy of nearby Type\,Ia supernovae have been presented in relation to the study of \ion{Na}{i} and \ion{Ca}{ii} to probe the circumstellar matter associated with the progenitors of these events. This has led to the discovery of time variability in high-resolution \ion{Na}{i} absorption line systems for some supernovae (see e.g. \citealt{2007Sci...317..924P}; \citealt{2009ApJ...702.1157S}). Several similar events were detected using low-resolution spectroscopy (\citealt{2009ApJ...693..207B}; \citealt{2010AJ....140.2036S}). However, not all Type\,Ia SNe show this behaviour (e.g. \citealt{2013A&A...549A..62P}). Furthermore, \citet{2011Sci...333..856S} found that for a large sample of Type\,Ia SNe, the absorbing material (\ion{Na}{i}) tends to be blue-shifted with respect to the strongest sodium absorption line, which has been suggested as an indication of outflows from the supernova progenitor systems. This is expected for single-degenerate (SD) systems but is much less probable for double-degenerate (DD) systems. This result needs to be taken with caution as the strongest sodium absorption component does not necessarily have the same velocity as the disk gas at the position of the supernova (e.g. \citealt{2012ApJ...748L..11R}). \citet{2013ApJ...779...38P} showed that some supernovae -- all of which were classified as blue-shifted -- display anomalously high \ion{Na}{i} column densities with respect to the amount of dust extinction, suggesting a relation to the outflowing circumstellar gas. Pathways for producing this enhancement were given for both the SD and DD scenarios. The increasing availability of high-resolution high-sensitivity spectra (in particular when averaging over multiple epochs) has provided the possibility to study the physical properties of the ISM in other galaxies in detail. Until recently, only a few supernovae, such as SN\,1987A, could be studied in detail. Now, bright stars in the Magellanic Clouds (\citealt{2006A&A...447..991C}; \citealt{2006ApJS..165..138W}) and the Andromeda and Triangulum galaxies (\citealt{2008A&A...480L..13C,2008A&A...492L...5C}) are accessible - with some effort - to optical (high-resolution) spectroscopy. To proceed in studying distant galaxies it is becoming readily possible to use nearby SNe, similar to the exploitation of SN\,1987A to study the LMC (\citealt{1987ESOC...26..511G}; \citealt{1988PASAu...7..527P}), as background candles to probe the (physical) properties of the ISM such as dust extinction and UV field strength in their host galaxy through the observation of atoms, molecules, and diffuse interstellar bands (DIBs). However, sufficiently bright SNe are required to obtain high-S/N spectra at high-resolution, which limits the number of possibilities for such studies. Furthermore, not all SNe will probe significant columns of interstellar matter, in particular if they are located in the approaching side of the respective galaxy. Although rare, an increasing number of SNe have been observed with intermediate to high spectral resolution to probe the atomic and molecular diffuse ISM, including DIBs, in the disks and halos of galaxies beyond the Local Group, such as the Centaurus group (\citealt{1985ApJ...299..852D}; \citealt{1987AJ.....94..651R}; \citealt{1989A&A...215...21D}), the Virgo cluster (\citealt{1990AJ.....99.1476S}; \citealt{1991ApJ...383L..41M}) and the M81 group (\citealt{1994ApJ...420L..71B}). The most detailed studies of both resolved molecular absorption lines and resolved DIB profiles arising from the cold diffuse ISM in galaxies beyond the Local Group were presented in \citet{2005A&A...429..559S} and \citet{2008A&A...485L...9C}. In addition, interstellar absorption features such as the UV bump and the strong 5780 DIB have also been detected in distant damped Ly$\alpha$ (DLA) systems (\citealt{2004ApJ...614..658J}; \citealt{2006ApJ...647L..29Y}; \citealt{2008MNRAS.383L..30E}, \citealt{2012MNRAS.tmp..159S}). Here we present a detailed study of the properties of the ISM in the SN host galaxy and the Galactic halo as probed by SN\,2008fp. The spectra were obtained as part of an ongoing campaign to observe atomic line variablity in bright Type\,Ia SNe in an effort to distinguish between the SD and DD scenarios for Type\,Ia SNe (cf. \citealt{2013A&A...549A..62P}). In Sect.~\ref{sec:obs} we present the observations and data reduction and give basic properties of the SN and its host galaxy. The detected Galactic and extra-galactic interstellar absorption lines are presented in Sect.~\ref{sec:sn2008fp-ISM}. In Sect.~\ref{sec:discussion} we discuss the molecular content, physical cloud conditions, interstellar line variability and the presence and behaviour of extra-galactic molecules and DIBs in the SN host galaxy. We conclude with a summary in Sect.~\ref{sec:conclusion}.	\label{sec:conclusion} We presented deep high-resolution optical spectra of the Type\,Ia supernovae 2008fp. We focused on the analysis and interpretation of these spectra in the context of understanding the physical properties of interstellar molecular clouds in extra-galactic environments. This analysis of the line-of-sight towards SN\,2008fp revealed the following: \begin{itemize} \item In addition to the main atomic species (\ion{Na}{i}, \ion{Ca}{ii}), many weaker lines are detected: \ion{Fe}{i}, \ion{Ca}{i}, \ion{Ti}{ii}, CH, CH$^+$, CN, as well as diffuse interstellar bands at 5780, 5797, 5849, 6196, and 6283~\AA. \item The 5780 and 5797~\AA\ DIBs indicate that the $\zeta$-type cloud component at 1770~\kms\ can be considered a translucent cloud sufficiently shielded from interstellar UV radiation (from young stars in the host galaxy) to ensure a rich molecular chemistry. From a comparision with Galactic $\sigma$ and $\zeta$-type DIB spectra we inferred \Ebv = 0.45~mag for the total extra-galactic reddening towards SN\,2008fp. This fully agrees with $A_\mathrm{V}$ and $R_\mathrm{V}$ derived from modelling the light curve. \item The molecular hydrogen fraction of 0.7$^{0.15}_{-0.2}$ corresponds to that of translucent clouds. \item The C$_2$ (2,0) band is detected beyond the Local Group, in the SN\,2008fp host galaxy. The rotational line analysis yielded T$_\mathrm{rot}~\approx$~30~K and $n_C~\approx$~250~cm$^{-3}$. The C$_2$/H$_\mathrm{total}$ fraction is a factor three or more higher than typically found for Galactic diffuse clouds. \item We presented a column density of N(C$_3$) = 1.7 $\pm$ 0.5 $\times$ 10$^{13}$~cm$^{-2}$ for extra-galactic C$_3$ towards SN\,2008fp, which constitutes the first detection of C$_3$ beyond the Local Group. The relative abundance of C$_3$ with respect to C$_2$ is similar to that of the Galactic mean. \item The relative depletion of Ti/H is $\approx$ -9.2~dex, which indicates a higher level of depletion than the -7.1~dex depletion found for the solar neighbourhood. \item From the atomic and molecular line analysis we conclude that the bulk of reddening (dust extinction), and therefore of polarisation, occurs within molecular clouds in the SN\,2008fp host galaxy along the line of sight. If not physically associated with the SN progenitor, this implies that the reddening and polarisation arise primarily in the ISM and not in the CSM. This conclusion is strengthened by the polarisation position angle. If the dust were of CS nature, it would produce polarisation by scattering, which would turn into a non-null net polarisation only if the dust is distributed asymmetrically. \item The wavelength dependency of interstellar polarisation in the host of SN\,2008fp clearly differs from that displayed by similarly reddened stars in the Galaxy. The observed behaviour is very similar to that of SN\,2006X, which also shows a total-to-selective extinction ratio $R_V<2$, indicating a dust size/composition significantly different from that typical of the Milky Way. \item From the CN(0,0) band the CN excitation temperature at 2.64~mm is 2.9$\pm$0.4~K, consistent with that of the cosmic microwave background, and statistically inconclusive regarding local collisional excitations by electrons. \item The lack of variability over a 28-day period excludes small, isolated clouds with sizes similar to the SN photospheric radius of 100 to 200~AU, supporting the presence of a more patchy - fractal - ISM. \item High-resolution spectra of Type\,Ia SNe reveal several sightlines -- SN\,2006X, SN\,2008fp, SN\,2009ig -- with an inexplicably strong CN B-X (0,0) absorption band. Whether and how this is related to a specific subset of Type\,Ia SNe or their host galaxies remains an open question. \end{itemize}	14	3	1403.4386	Context. Supernovae (SNe) offer a unique opportunity to study physical properties, small-scale structure, and complex organic chemistry of the interstellar medium (ISM) in different galaxies. <BR /> Aims: Interstellar absorption features, such as atomic and molecular lines as well as diffuse interstellar bands (DIBs), can be used to study the physical properties of extra-galactic diffuse interstellar clouds. <BR /> Methods: We used optical high-resolution spectroscopy to study the properties of the ISM in the SN 2008fp host galaxy, ESO 428-G14. The properties of intervening dust were investigated via spectropolarimetry. <BR /> Results: The spectra of SN 2008fp reveal a complex of diffuse atomic clouds at radial velocities in line with the systematic velocities of the host galaxy. In addition, a translucent (A<SUB>V</SUB> ~ 1.5 mag) cloud is detected at a heliocentric velocity of 1770 km s<SUP>-1</SUP> (redshifted by 70 km s<SUP>-1</SUP> with respect to the system velocity). This cold dense cloud is rich in dense atomic gas tracers, molecules, as well as DIBs. We have detected both C<SUB>2</SUB> and C<SUB>3</SUB> for the first time in a galaxy beyond the Local Group. The CN (0, 0) band-line ratios are consistent with an excitation temperature of T = 2.9 ± 0.4 K. The interstellar polarisation law deviates significantly from what is observed in the Galaxy, indicating substantial differences in the host dust/size composition. No variations over a period of about one month are observed in any of the ISM tracers. <BR /> Conclusions: The lack of variability in the extra-galactic absorption line profiles implies that the absorbing material is not circumstellar and thus not directly affected by the SN event. It also shows that there are no significant density variation in the small-scale structure of the molecular cloud down to 100 AU. C<SUB>2</SUB> is used to probe the cold diffuse ISM density and temperature. Here we also use observations of CN in a distant galaxies, though for now still in a limited way, for in situ measurements of the cosmic background radiation temperature. <P />This paper is based on observations made with ESO Telescopes at Paranal Observatory under program IDs 081.D-0558, 081.D-0697, and 082.D-0004.	false	[ "extra-galactic diffuse interstellar clouds", "radial velocities", "diffuse atomic clouds", "different galaxies", "temperature", "physical properties", "dense atomic gas tracers", "SN 2008fp", "diffuse interstellar bands", "the SN 2008fp host galaxy", "DIBs", "line", "situ measurements", "ESO Telescopes", "082.D-0004", "program IDs", "the cosmic background radiation temperature", "SN", "ISM", "complex organic chemistry" ]	3.439641	4.970264	-1
745783	[ "Molinari, D.", "Gruppuso, A.", "Polenta, G.", "Burigana, C.", "De Rosa, A.", "Natoli, P.", "Finelli, F.", "Paci, F." ]	2014MNRAS.440..957M	[ "A comparison of CMB angular power spectrum estimators at large scales: the TT case" ]	16	[ "Dipartimento di Fisica e Astronomia, Università degli Studi di Bologna, viale Berti Pichat 6/2, I-40127 Bologna, Italy; INAF-IASF Bologna, Via Piero Gobetti 101, I-40129 Bologna, Italy", "INAF-IASF Bologna, Via Piero Gobetti 101, I-40129 Bologna, Italy; INFN, Sezione di Bologna, Via Irnerio 46, I-40126 Bologna, Italy", "Agenzia Spaziale Italiana Science Data Center, c/o ESRIN, via Galileo Galilei, I-00044 Frascati, Italy; INAF - Osservatorio Astronomico di Roma, via di Frascati 33, I-00040 Monte Porzio Catone, Italy", "INAF-IASF Bologna, Via Piero Gobetti 101, I-40129 Bologna, Italy; Dipartimento di Fisica e Scienze della Terra, Università degli Studi di Ferrara, via Giuseppe Saragat 1, I-44100 Ferrara, Italy", "INAF-IASF Bologna, Via Piero Gobetti 101, I-40129 Bologna, Italy", "INAF-IASF Bologna, Via Piero Gobetti 101, I-40129 Bologna, Italy; Agenzia Spaziale Italiana Science Data Center, c/o ESRIN, via Galileo Galilei, I-00044 Frascati, Italy; Dipartimento di Fisica e Scienze della Terra, Università degli Studi di Ferrara, via Giuseppe Saragat 1, I-44100 Ferrara, Italy; INFN, Sezione di Ferrara, via Giuseppe Saragat 1, I-44100 Ferrara, Italy", "INAF-IASF Bologna, Via Piero Gobetti 101, I-40129 Bologna, Italy; INFN, Sezione di Bologna, Via Irnerio 46, I-40126 Bologna, Italy", "SISSA - Scuola Internazionale Superiore di Studi Avanzati, via Bonomea 265, I-34136 Trieste, Italy" ]	[ "2015IJMPD..2444008G", "2015MNRAS.448.2854C", "2016A&A...593A..15D", "2016A&A...594A...2P", "2017MNRAS.466.3961S", "2017PhRvD..95d3532G", "2018A&A...609A..52B", "2018PhRvD..98b3521M", "2018PhRvD..98j3526V", "2019JCAP...10..001L", "2019JCAP...12..052N", "2019PhRvD.100b3538L", "2020A&A...641A...7P", "2020JCAP...11..066G", "2021JCAP...01..026L", "2021JCAP...07..034B" ]	[ "astronomy" ]	6	[ "methods: data analysis", "methods: numerical", "methods: statistical", "cosmic background radiation", "cosmology: observations", "cosmology: theory", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1967ApJ...147...73S", "1973ApJ...185..757H", "1980PhLB...91...99S", "1981PhRvD..23..347G", "1982PhLB..108..389L", "1982PhRvL..48.1220A", "1985SvAL...11..271K", "1994PhRvL..73.3347K", "1996ApJ...473..576F", "1997PhRvD..55.5895T", "1999ApJ...512..511M", "2001PhRvD..64f3001T", "2002ApJ...567....2H", "2003ApJS..148..135H", "2003ApJS..148..233P", "2004ApJ...600...32K", "2004ApJ...609....1J", "2004ApJS..155..227E", "2004MNRAS.348..885E", "2004MNRAS.349..603E", "2004PhRvD..70h3511W", "2005ApJ...622..759G", "2005JCAP...11..001P", "2005MNRAS.358..833T", "2006ApJ...645L..89S", "2006ApJ...647..823J", "2007ApJ...665...55W", "2008MNRAS.387..209M", "2008MNRAS.389.1284T", "2009ApJ...692.1247P", "2009ApJS..180..306D", "2009MNRAS.400..219B", "2009MNRAS.400..463G", "2009PhRvD..79l3515G", "2010ApJ...714L.265K", "2010MNRAS.407..399P", "2011ApJS..192...17B", "2011MNRAS.411.1445G", "2011MNRAS.412.2383C", "2013ApJS..208...19H", "2013ApJS..208...20B", "2013JCAP...07..047G", "2013MNRAS.434.3071P", "2014A&A...571A...1P", "2014A&A...571A...2P", "2014A&A...571A..15P", "2014A&A...571A..16P", "2014A&A...571A..20P", "2014A&A...571A..23P", "2014A&A...571A..24P" ]	[ "10.1093/mnras/stu386", "10.48550/arXiv.1403.1089" ]	1403	1403.1089_arXiv.txt	The pattern of the cosmic microwave background (CMB) anisotropy field can be used to probe cosmology to high precision, as shown by the Wilkinson Microwave Anisotropy Probe (WMAP) 9 years results \citep{Hinshaw:2012fq} and by the very recent {\it Planck} cosmological results (see \cite{Ade:2013xsa} and references therein). CMB data have given a significant contribution in setting up the $\Lambda$ cold dark matter ($\Lambda$CDM) cosmological concordance model. The latter establishes a set of basic quantities for which CMB observations and other cosmological and astrophysical data-sets agree\footnote{See, however, \cite{Ade:2013lmv} for a possible tension concerning the $\Omega_m$ estimate from {\it Planck} CMB and galaxy clusters data.}: spatial curvature close to zero; $\sim 68.5 \%$ of the cosmic density in the form of Dark Energy; $\sim 26.5\%$ in cold dark matter; $\sim 5\%$ in baryonic matter; and non perfectly scale invariant adiabatic, primordial perturbations compatible with Gaussianity \citep{Ade:2013lta,Ade:2013ydc}. In particular, the largest scales of the temperature anisotropies map are of great interest because they directly probe the Early Universe (or Inflationary Phase of the Universe \cite{Starobinsky:1980te,Guth:1980zm,Linde:1981mu,Albrecht:1982wi}). They correspond to angular scales larger than the horizon size at decoupling as observed today, i.e. $\theta > 2^{\circ}$ or, equivalently $\ell < \ell_{dec} \sim 90$ (see for example \cite{Page2003}) with $\ell $ being the multipole order of the spherical harmonics expansion \begin{equation} {\delta T}(\hat n) = \sum_{\ell m} a_{\ell m} \, Y_{\ell m} (\hat n) \label{SHE} \, , \end{equation} where $\delta T = T(\hat n) - T_0$ is the temperature anisotropy observed in the direction $\hat n$ relative to the CMB average temperature $T_0 \simeq 2.725 K$ (\cite{Mather1999}; see also \cite{Fixsen1996} for a constraint of the CMB black body shape), and with $a_{\ell m}$ being the coefficients of the Spherical Harmonics $Y_{\ell m} (\hat n) $. The main contribution to the CMB anisotropies at these large scales is provided by the so called Sachs-Wolfe effect and by a subdominant integrated Sachs-Wolfe effect \citep{Sachs:1967er} which is different from zero because of the recent (from a cosmological point of view) transition to an accelerated phase of the Universe \citep{Kofman:1985fp} likely associated to a dark energy (or Cosmological constant) component. In principle a stochastic background of primordial gravitational waves can also give a contribution to the temperature CMB anisotropies at these largest scales, depending on the tensor-to-scalar ratio, $r$, constrained by the current data \citep{Hinshaw:2012fq,Ade:2013lta}. A firm detection of the primordial gravitational waves requires CMB B modes polarization measurement at large scale \citep{Knox:1994qj}. From the observational point of view the CMB anisotropies temperature map, as observed by WMAP 9 year, is cosmic variance dominated, i.e., cosmic variance exceeds the instrument noise, up to $\ell = 946$, see \citep{Bennett:2012fp}. For {\it Planck} data this crossing happens at $\ell \sim 1500$ \citep{Planck:2013kta}. Therefore, at the largest scales ($\ell < \ell_{dec}$), the effect of instrumental noise is almost negligible. Keeping this in mind, it is even more important to employ the most accurate data analysis tools. In the current paper we focus on the angular power spectrum (APS), which is the main observable for diagnosis of the CMB map. The method that is capable to provide APS with no bias and with the minimum variance, as provided by the Fisher-Cramer-Rao inequality is the Quadratic Maximum Likelihood (QML) method \citep{tegmark_tt,tegmark_pol}. Such an optimal method has the drawback of being computationally expensive and then limited by the number of pixels. It is currently implemented and applied at low resolution (see e.g. \cite{Gruppuso2009}). Several other strategies for measuring $C_{\ell}$ at low resolution have been developed and applied to CMB data with excellent results. These methods include different sampling techniques such as Gibbs \citep{jewell, wandelt, eriksen}, adaptive importance \citep{teasing} and Hamiltonian \citep{taylor}. At high multipoles ($\ell > 30$, \cite{efstathiou}) the so called pseudo-$C_{\ell}$ algorithms are usually preferred to others techniques. These methods, in fact, implement the estimation of power spectral densities from periodograms \citep{hauser}. Basically, they estimate the $C_l$ through the inverse Harmonical transform of a masked map that is then deconvolved with geometrical kernels and corrected with a noise bias term. These techniques, such as Master \citep{master}, Cross-Spectra \citep{saha,polenta,grain}, give unbiased estimates of the CMB power spectra and moreover, it has been shown they work successfully when applied to real data at high multipoles \citep{ACBAR,Jones:2006,MAXIPOL,dunkley_wmap5}. These estimators are pretty quick and light from a computational point of view. However, it is well known that at low multipoles they are not optimal since they provide power spectra estimates with error bars larger than the minimum variance. We note that in \cite{efstathiou2} an hybrid approach has already been proposed, i.e. the QML at low and the pseudo-$C_{\ell}$ at high multipoles, where a recipe for an hybrid covariance is consistently given. However, in that paper the QML is considered only up to ${\ell}=40$ since it is applied to maps with a full width at half maximum (FWHM) $=3 ^{\circ}$. The aim of the present work is to compare quantitatively the QML and pseudo-$C_{\ell}$ methods under {\it realistic} assumptions focusing on the largest scales (i.e. low multipoles) where CMB anomalies are mainly located \citep{anomalies,Ade:2013sta}. The idea is to provide the scientific community with a reference analysis comparing the heavy QML method and the light and quick pseudo-$C_{\ell}$ approach in the range of multipoles $\ell < \ell_{dec} \sim 90$. Here we quantitatively address this problem through realistic Monte Carlo (MC) simulations. A similar analysis has been carried out by WMAP team using the so-called $C^{-1}$ method employed in \cite{Bennett:2012fp} that shows an improvement with respect to the pseudo-$C_{\ell}$ method mostly located at very low multipoles and in the intermediate S/N regime. Moreover, we evaluate the benefit of considering an optimal APS estimation taking into account some examples of estimators of anomalies that are commonly used at large scales to test the consistency of the observations with the $\Lambda CDM$ model. The paper is organized as follows. In Section \ref{descriptionscodes} we describe the considered implementations of QML and pseudo-$C_{\ell}$ codes, called respectively {\it BolPol} and {\it cROMAster}. In Section \ref{computationalrequirements} the computational requirements for both the methods are given. Section \ref{comparison} is devoted to the detailed description of the comparison and corresponding results. Examples on how the previous analysis propagates to the APS based estimators, commonly used in literature for the analysis of the large scales anomalies, are provided in Section \ref{WMAP}. Conclusions are drawn in Section \ref{conclusions}.	\label{conclusions} Our main result is given by Fig. \ref{variancemaskedsky} and \ref{varianzarelativamaskedsky} where the intrinsic variance, see equation (\ref{figureofmeritdef2}), of the two APS estimators, namely {\it BolPol} and {\it cROMAster}, are compared under realistic conditions. We have found that the QML method is markedly preferable in the range $2 \leq \ell \leq 100$. Moreover, we note that the largest difference between the two codes is for the lowest multipoles: for $\ell$ smaller than $\sim 20$ the square root of the intrinsic variance introduced in the estimates by the pseudo-$C_{\ell}$ is {\it at least} up to three times (two times) the QML one when we consider the WMAP kq85 mask (respectively the kq85 mask enlarged by 8 degrees). For higher multipoles (i.e. $\ell \geq 100$) we observe an opposite behaviour. This stems from the smoothing of the input maps that in turn it is a consequence of the adopted resolutions. Note however that when the two codes are run at the same resolution, i.e. $N_{side} = 64$, the QML has always a smaller variance than cROMAster in the commonly valid multipole domain, as shown in Fig. \ref{variancemaskedskylowres}. We have also analysed how the intrinsic variance of the two APS methods impacts on some typical large scales anomaly estimators like the TT Parity estimator and Variance estimator. In conclusion, the use of {\it BolPol} for low resolution map analysis will bring to tighter constraints for these kind of estimators. Therefore we suggest to use the QML estimator and not the pseudo-$C_{\ell}$ method in order to perform accurate analyses that are based on the APS at large angular scales (at least $\ell \leq \ell_{dec} \simeq 90$). This might be of particular interest for studying large scale anomalies in the temperature anisotropy pattern. In a future work we will extend this analysis to the polarization field, which is crucial to reveal the reionization imprints at large scale.	14	3	1403.1089	In the context of cosmic microwave background (CMB) data analysis, we compare the efficiency at large scale of two angular power spectrum algorithms, implementing, respectively, the quadratic maximum likelihood (QML) estimator and the pseudo-spectrum (pseudo-C<SUB>ℓ</SUB>) estimator. By exploiting 1000 realistic Monte Carlo simulations, we find that the QML approach is markedly superior in the range 2 ≤ ℓ ≤ 100. At the largest angular scales, e.g. ℓ ≤ 10, the variance of the QML is almost 1/3 (1/2) that of the pseudo-C<SUB>ℓ</SUB>, when we consider the WMAP kq85 (kq85 enlarged by 8°) mask, making the pseudo-spectrum estimator a very poor option. Even at multipoles 20 ≤ ℓ ≤ 60, where pseudo-C<SUB>ℓ</SUB> methods are traditionally used to feed the CMB likelihood algorithms, we find an efficiency loss of about 20 per cent, when we considered the WMAP kq85 mask, and of about 15 per cent for the kq85 mask enlarged by 8°. This should be taken into account when claiming accurate results based on pseudo-C<SUB>ℓ</SUB> methods. Some examples concerning typical large-scale estimators are provided.	false	[ "kq85", "the pseudo-spectrum estimator", "the WMAP kq85 mask", "large scale", "pseudo-C<SUB>ℓ</SUB> methods", "QML", "WMAP", "CMB", "the kq85 mask", "typical large-scale estimators", "the WMAP kq85", "spectrum (pseudo-C<SUB>ℓ</SUB>) estimator", "C<SUB>ℓ</SUB", "the pseudo", "2 ≤ ℓ ≤", "20 ≤ ℓ ≤", "≤", "the CMB likelihood", "the quadratic maximum likelihood", "the largest angular scales" ]	12.649879	1.803792	103
12398957	[ "Li, Huiquan", "Yu, Cong", "Wang, Jiancheng", "Xu, Zhaoyi" ]	2015PhRvD..92b3009L	[ "Force-free magnetosphere on near-horizon geometry of near-extreme Kerr black holes" ]	9	[ "Yunnan Observatories, Chinese Academy of Sciences, 650011 Kunming, China and Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, 650011 Kunming, China", "Yunnan Observatories, Chinese Academy of Sciences, 650011 Kunming, China and Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, 650011 Kunming, China", "Yunnan Observatories, Chinese Academy of Sciences, 650011 Kunming, China and Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, 650011 Kunming, China", "Yunnan Observatories, Chinese Academy of Sciences, 650011 Kunming, China and Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, 650011 Kunming, China" ]	[ "2012arXiv1203.3561C", "2014PhRvD..90l4009Z", "2015LNP...906.....B", "2015PhRvD..92b4049Z", "2015arXiv150504781W", "2017LRR....20....1C", "2017MNRAS.468.4351C", "2017PhRvD..96b3014L", "2018PhRvD..97f4017L" ]	[ "astronomy", "physics" ]	5	[ "97.60.Lf", "11.25.Tq", "Black holes", "Gauge/string duality", "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena", "High Energy Physics - Theory" ]	[ "1977MNRAS.179..433B", "1977MNRAS.179..457Z", "1999PhRvD..60j4030B", "2005ApJ...635.1197M", "2007GReGr..39..785M", "2008PhRvD..78b4004T", "2009JHEP...12..037C", "2009PhRvD..80l4008G", "2011MNRAS.417.1098M", "2011NuPhS.216..194B", "2013CQGra..30s5012B", "2014JHEP...12..185L", "2014MNRAS.445.2500G", "2014PhRvD..89j6011W", "2014PhRvD..90d4038P", "2014PhRvD..90f4045H", "2015JHEP...07..090L" ]	[ "10.1103/PhysRevD.92.023009", "10.48550/arXiv.1403.6959" ]	1403	1403.6959_arXiv.txt	\label{sec:introduction} Black hole magnetospheres as well as many high-energy astrophysical objects, for instance, relativistic jets in active galactic nuclei and gamma-ray bursts and probably ultra-strongly magnetized neutron stars (magnetars) involve relativistically magnetically dominated plasma. Under such circumstances, magnetic fields play crucial roles in the dynamics of these astrophysical scenarios, which can drive powerful winds/jets from these astrophysical objects. In addition, the magnetic dissipation can also give rise to remarkable non-thermal emissions in the these high-energy astrophysical phenomena. It is widely accepted that, in these objects, the magnetic energy density conspicuously exceeds the thermal and rest mass energy density of particles. The force-free electrodynamics behave well in such extreme magnetically dominated scenarios as the less important terms, such as the inertia and pressure, are entirely ignored. In the framework of force-free electrodynamics, Blandford and Znajeck (BZ) \cite{Blandford:1977ds} proposed that the rotation energy of a kerr black hole could be extracted via the action of force free electromagnetic fields, in the form of Poynting flux via magnetic field lines penetrating the central black hole. The tapping of the rotational energy form spinning black hole via magnetized plasma appears to be the most astrophysically promising mechanism to exploit the black hole energy. Relativistic jets in active galactic nuclei, galactic microquasars, and gamma ray bursts may well be driven by the BZ mechanism. Exploring solutions in this configuration is crucial for examining the energy extraction process and checking the efficiency of this mechanism. However, very few analytical solutions have been obtained so far (a collection of available solutions is summarised in \cite{Gralla:2014yja} with modern treatments). In particular, it is hard to derive the solutions in the case of rapidly rotating holes. A key reason is that little information can be extracted from the single stream equation, accompanied with a regularity condition on the horizon \cite{Znajek:1977unknown}. In the original work \cite{Blandford:1977ds}, split monopole and paraboloidal solutions for slowly rotating black holes are obtained with expansion on the angular momentum of the hole to the leading order. But, this perturbative approach involving higher order corrections seems unsuccessful in searching for solutions for rapidly rotating back holes \cite{Tanabe:2008wm}. Nevertheless, a set of exact solutions for arbitrary angular momentum are obtained in the far-field region by Menon and Dermer (MD solutions) \cite{Menon:2005va,Menon:2005mg}, which indeed implies the formation of collimated outflow or jet \cite{Menon:2011zu}. Based on previous solutions, generlised solutions to time-dependent and non-axisymmetric case are obtained by assuming the four-current to be along one of the null tetrads \cite{Brennan:2013jla}. In theoretical physics, extreme and near-extreme black holes attract plenty of attention in past years, due to the proposal of the AdS/CFT correspondence. For extremely rapidly rotating black holes, this specific dual description is the Kerr/CFT correspondence \cite{Guica:2008mu,Castro:2009jf,Bredberg:2011hp}, constructed on the near-horizon geometry of (near-)extreme Kerr black holes, namely (near-)NHEK \cite{Bardeen:1999px,Castro:2009jf}. In this paper, we shall adopt the near-NHEK, replacing the full Kerr metric, to study the force-free magnetosphere around near-extreme Kerr black holes. The near-NHEK should be very suitable for seeking for solutions in this system. First, it has a very simple form with precise symmetries, so that we can simplify the equations and study them analytically. Second, in the near-horizon region, we can take full advantage of the constraint condition on the horizon. Hopefully, the solution obtained on near-NHEK may provide us clues for searching for solutions in the full Kerr background. Meanwhile, the astrophysical processes, including the BZ process near rapidly rotating holes, may also provide an arena for testing the Kerr/CFT correspondence. Recently, there appear the works \cite{Porfyriadis:2014fja,Hadar:2014dpa} along this line, where the authors studied the gravity waves produced from massive objects orbiting around an (near-)extreme Kerr on (near-)NHEK in the context of Kerr/CFT. In \cite{Wang:2014vza}, the dual CFT aspects of the magnetosphere around slow-rotating balck holes on AdS space has been discussed. Here, we try to adopt the near-NHEK to explore analytical solutions of force-free magnetospheres near rapidly rotating holes. The paper is organised as follows. In Section.\ \ref{sec:adsdescription}, we present the near-NHEK and the expressions of EM fields on it. In Section.\ \ref{sec:freemag}, we give the stream equation of force-free magnetospheres on near-NHEK and then seek for solutions to the equation. In Section.\ \ref{sec:extraction}, we show that it is possible to form a outward flux from the derived solution. A brief summary is made in the last section. During the proceedings of this work, the works \cite{Lupsasca:2014pfa,Lupsasca:2014hua} appear, which focus on magnetospheres on exact NHEK by taking advantage of the conformal symmetries of the geometry.	\label{sec:conclusion} It is possible that the (near-)NHEK makes a good tool to analytically study the force-free magnetosphere near rapidly rotating astronomical black holes. By focusing on the near-horizon region, we showed that the boundary condition on the horizon can be derived directly from the stream equation on near-NHEK. We got consistent solutions near the rotation axis which imply the formation of outflow. The angular dependence of the solutions resembles that of the asymptotic solutions obtained in \cite{Menon:2005va,Menon:2005mg,Menon:2011zu}.	14	3	1403.6959	We study force-free magnetospheres in the Blandford-Znajek process from rapidly rotating black holes by adopting the near horizon geometry of near-extreme Kerr black holes (near-NHEK). It is shown that the Znajek regularity condition on the horizon can be directly derived from the resulting stream equation. In terms of the condition, we split the full stream equation into two separate equations. Approximate solutions around the rotation axis are derived. They are found to be consistent with previous solutions obtained in the asymptotic region. The solutions indicate energy and angular-momentum extraction from the hole.	false	[ "two separate equations", "the resulting stream equation", "black holes", "the full stream equation", "near-extreme Kerr black holes", "the near horizon geometry", "Kerr", "previous solutions", "Znajek", "Approximate solutions", "the hole", "the Znajek regularity condition", "the asymptotic region", "the horizon", "-NHEK", "energy and angular-momentum extraction", "force-free magnetospheres", "the Blandford-Znajek process", "terms", "the rotation axis" ]	8.111993	4.428046	28
745674	[ "Walker-Smith, S. L.", "Richer, J. S.", "Buckle, J. V.", "Hatchell, J.", "Drabek-Maunder, E." ]	2014MNRAS.440.3568W	[ "The James Clerk Maxwell Telescope dense gas survey of the Perseus molecular cloud" ]	13	[ "Astrophysics Group, Cavendish Laboratory, J J Thomson Avenue, Cambridge CB3 0HE, UK", "Astrophysics Group, Cavendish Laboratory, J J Thomson Avenue, Cambridge CB3 0HE, UK; Kavli Institute for Cosmology, Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Astrophysics Group, Cavendish Laboratory, J J Thomson Avenue, Cambridge CB3 0HE, UK; Kavli Institute for Cosmology, Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "School of Physics, University of Exeter, Stocker Rd, Exeter EX4 4QL, UK", "School of Physics, University of Exeter, Stocker Rd, Exeter EX4 4QL, UK; Blackett Laboratory, Imperial College London, Prince Consort Rd, London SW7 2AZ, UK" ]	[ "2015ApJ...798...61T", "2016A&A...586A..44C", "2016ApJ...823..151C", "2017RSOS....470754R", "2018A&A...620A..30M", "2018ApJ...862...42T", "2018MNRAS.477.2455C", "2019A&A...632A.101T", "2020ApJ...891...66N", "2021A&A...655A..65T", "2021A&A...656A.146M", "2021ApJS..254...14Y", "2023ApJ...942...32Y" ]	[ "astronomy" ]	2	[ "stars: formation", "ISM: clouds", "ISM: individual objects: Perseus", "ISM: jets and outflows", "ISM: molecules", "Astrophysics - Astrophysics of Galaxies" ]	[ "1976ApJ...209L.137Z", "1980ApJ...239L..17S", "1982ApJ...262..590F", "1987ARA&A..25...23S", "1987MNRAS.226..461J", "1990ApJ...348..530C", "1990MNRAS.244..668P", "1991MNRAS.252..442P", "1992A&A...261..274C", "1996A&A...311..858B", "1997ApJ...484..256G", "1997IAUS..182..153B", "1998A&A...335..266B", "1998ApJ...499..777B", "1999MNRAS.305..651F", "2000AJ....120.1467W", "2000ApJ...536..857H", "2000ApJ...537..245H", "2001ApJ...546L..49S", "2001ApJ...555..139C", "2002A&A...383..892B", "2004A&A...413..993J", "2004A&A...415.1021J", "2004A&A...416..603J", "2004A&A...426..503W", "2004MNRAS.351.1054R", "2005A&A...432..369S", "2005A&A...440..151H", "2006ApJ...636L.137P", "2007A&A...468..627V", "2007A&A...468.1009H", "2007A&A...472..187H", "2007prpl.conf..245A", "2008MNRAS.387..954D", "2009A&A...502..199B", "2009MNRAS.399.1026B", "2010A&A...522A..91T", "2010MNRAS.406.2488D", "2010MNRAS.408.1516C", "2011A&A...535A..44F", "2011BASI...39..289G", "2011IAUS..270..451P", "2011IAUS..280...88T", "2011PASJ...63....1H", "2013A&A...549A..16N", "2013A&A...556A..76V", "2013MNRAS.431.1719M", "2014A&A...563L...3C" ]	[ "10.1093/mnras/stu512", "10.48550/arXiv.1403.3051" ]	1403	1403.3051_arXiv.txt	\subsection{Outflows and chemistry} High-velocity molecular outflows were first discovered in 1976 \citep{1976ApJ...209L.137Z}, and are thought to be one of the earliest observable signatures of star formation --- nearly all core-collapse candidate protostars are found to have outflows. Indeed outflows are thought to be integral to the star formation process, as they remove excess angular momentum \citep{2007prpl.conf..245A}. However despite the ubiquity of outflows, there is still much that we do not understand about them. One example is the outflow driving mechanism: they are thought to be powered by the gravitational energy of the infalling material in a contracting core \citep{1980ApJ...239L..17S}. However, as their source is usually deeply embedded in a dense infalling envelope, it is difficult to observe directly the outflow region and investigate the outflow driving mechanism. Molecular outflows are thought to develop from the entrainment of low-velocity flows by well-collimated (optical) jets --- the primary jet injects its momentum into the surrounding gas, resulting in a molecular outflow that traces the interaction between the jet and the surrounding environment \citep{1987ARA&A..25...23S}. Wide-angled winds sweeping up a shell of gas have also been invoked to explain outflows such as HH~211 \citep{2006ApJ...636L.137P}. In these outflow models, emission from molecular outflows originates from ambient gas that has been accelerated by a supersonic wind, and has been shock-processed. Observations support this as outflows tend to be associated with other objects such as H{\sc ii} regions, HH objects, H$_2$ jets and H$_2$O masers \citep{2004A&A...426..503W}. Outflows and shock chemistry are therefore inextricably linked. Some outflows show a greater degree of molecular richness than others --- exciting a wide range of molecules and enhancing the abundance of some to many orders of magnitude over their standard molecular cloud abundances. These `chemically-active' outflows are not necessarily distinguishable using CO emission (although they do tend to be young Class 0 sources), and their activity is also highly transient \citep{2011IAUS..280...88T}. Although SiO is the canonical outflow tracer, molecules such as \HCOplus\ and HCN are increasingly used to study outflows: the higher transitions (e.g. $J = 4\to3$) trace both the dense gas associated with the protostellar objects, as well as any outflowing/infalling gas \citep{1997ApJ...484..256G}. This allows a comparison of the physical and chemical properties in the source envelopes, with those in the outflow lobes further away from the central sources. In addition, both these molecules have similar critical densities and excitation energies (see Table \ref{mol_trans}), and are found to have highly-correlated spatial distributions. However, recent studies have shown that the abundances of the two molecules are enhanced to different degrees in particular outflows --- \citet{2010A&A...522A..91T} have shown that the HCN abundance is enhanced by $\sim 10-100$ times more than the \HCOplus\ abundance in the outflow linewings. One theory by \citet{2000ApJ...536..857H} is that the differences between HCN and \HCOplus\ in the linewings can be attributed to the effect of magnetic field lines that are not aligned with the turbulent flow direction. The ions are trapped on the magnetic field lines while the neutrals follow the turbulent flow, resulting in the ionic species having narrower lines and suppressed linewings compared to the neutral species \citep{2000ApJ...537..245H}. This theory assumes that the two molecules coexist and sample the same parts of the molecular cloud, exposing them to the same dynamical processes. This is not necessarily the case, as \HCOplus\ and HCN have similar critical densities and excitation conditions but follow different chemical networks \citep{1990MNRAS.244..668P}, making it likely that differences in abundances between the two molecules are due to chemical effects. This difference in enhancements between the two molecules is therefore more commonly attributed to shock-chemistry, as HCN is known to be enhanced in shocks \citep{2004A&A...413..993J}. \citet{2011IAUS..280...88T} have suggested that there exists an incomplete knowledge of outflows and their chemistry due to a too-focused approach on a few selected objects, and a lack of a statistically-significant number of outflows explored using molecules other than CO. This dense gas survey of \HCOplus\ and HCN $J = 4\to3$ transitions, over multiple sub-regions in the Perseus molecular cloud, is thus very timely. These sub-regions have previously been observed in \twelveCO/\thirteenCO/\CeighteenO\ by \citet{2010MNRAS.408.1516C}, tracing the large-scale structure and kinematics of Perseus, as well as the protostellar activity occurring there. We have performed follow-up observations using \HCOplus\ and HCN ($J = 4\to3$ transitions) to characterise the denser ($n$ \textgreater $10^6$~\cmcube) gas present, and to investigate a large number of previously-mapped outflows using our molecules for comparison. \subsection{The Perseus molecular cloud} The Perseus molecular cloud is a well-studied region of low to intermediate-mass star formation. Distance estimates for Perseus range from 230 to 350~pc \citep{2011A&A...535A..44F}, \changed{and \citet {2011PASJ...63....1H} have recently calculated the distance to Perseus to be 235~pc by astrometry}; but for the purposes of this paper we assume the value of 250~pc used by \citet{2010MNRAS.408.1516C}, with whom we will be comparing our results. We investigate four sub-regions in the Perseus molecular cloud with differing degrees of turbulence, clustering and star formation activity. A comparison of these sub-regions allows us to investigate the effects of the environment on the dense gas structure and outflow properties. The four sub-regions are: \begin{enumerate} \item NGC~1333: This is thought to be the most active and clustered region of star formation in Perseus, and contains a large concentration of TT stars, HH objects, bipolar jets and outflows, all of which will affect the outflow properties that we calculate. It is fairly young --- possessing cores at an age of $\sim 1$~Myr \citep{2005A&A...440..151H}; it also has a lumpy and filamentary structure --- dust ridges extend between clumps of YSOs, and there are several cavities filled with high-velocity outflow gas. \citet{2001ApJ...546L..49S} found that the strongest submillimetre emission originates in the south, and is associated with the YSOs IRAS~2, 3 and 4. \item IC~348: This is a slightly older region than NGC~1333, but is still undergoing active star formation. The most well-known feature in the region is HH~211 --- a highly-collimated bipolar outflow driven by a Class 0 protostar, which has been the subject of several interferometric studies (e.g. Chandler \& Richer, 2001). \item L~1448: This region contains a large number of young Class 0 protostars, and has been found to be dominated by outflow activity, which argues for it being relatively young. In particular, the outflow L~1448-C has been studied in great detail \citep{2013A&A...549A..16N,2000AJ....120.1467W}. \item L~1455: This is the smallest and faintest of the four regions. It has the highest proportion of Class I sources (compared with Class 0s), which points to it being older than L~1448. There are still quite a few prominent outflows, several of which have H$_2$ objects associated with them \citep{2008MNRAS.387..954D}. \end{enumerate} \subsection{Outline} We present \HCOplus\ and HCN molecular data that we observed in the Perseus molecular cloud. Section \ref{sec:obs} presents an overview of the observations and data reduction procedure, and Section \ref{sec:structure} compares the spatial and velocity structures of the \HCOplus\ and HCN emission. We present an analysis of outflow properties in the 4 Perseus sub-regions, in both \HCOplus\ (Section \ref{sec:hcooutflows}) and in HCN (Section \ref{sec:hcnoutflows}), comparing the results obtained for the two molecules. Section \ref{sec:abundances} presents a comparison of the relative abundances of \HCOplus/HCN in IRAS~2 (NGC~1333) to investigate the chemistry within the outflow.	\label{sec:concl} We have carried out an analysis and comparison of the \HCOplus\ and HCN $J = 4\to3$ emission in several sub-regions of the Perseus molecular cloud and we have found the following results: \begin{enumerate} \item The \HCOplus\ shows much more extended structure than the HCN which is largely confined to compact clumps. We can constrain densities of the filamentary structures that are solely traced by \HCOplus\ to lie between the critical densities of the \HCOplus\ and HCN transitions. \item HCN is predominantly excited around protostellar objects, unlike the \HCOplus\ emission, which is also associated with a high proportion of starless SCUBA dust cores -- we can use the combination of a $3 \sigma$ detection for both molecules to pinpoint the truly protostellar cores. \item \HCOplus\ shows large-scale velocity structure across the individual sub-regions, particularly in NGC~1333; HCN on the other hand shows very little change in velocity across a particular sub-region, except around those protostars that power the most energetic outflows. \item \HCOplus\ emission is a good tracer of outflow activity -- it identifies over 50\% of the outflow sources that have been previously identified by \twelveCO, and those it misses are generally significantly weaker. \item \HCOplus\ outflow driving forces exhibit similar trends with $M_{\rm env}$ and $L_{\rm bol}$ when compared with \twelveCO, which indicates that they are excited similarly within the outflow. It also exhibits a similar trend with $M_{\rm env}$ to that calculated in Ophiuchus, indicating similar outflow driving mechanisms in the two separate regions. \item The outflow properties calculated for HCN are on average within a factor of 2 of those calculated for the same outflows using \HCOplus. This is an indication that the \HCOplus\ and HCN are excited similarly in the majority of the outflows that do not exhibit strong shock activity. \item HCN traces the most energetic outflows, especially in certain cases (e.g. IRAS~4) where the outflow wings extend further in velocity than in \HCOplus. The presence of large enhancements of HCN in particular outflows is thought to be an indication of shock activity within the outflow, and illustrates the youth and energy of the driving source. \item The increased enhancement of HCN in the blue lobe of outflow A in IRAS~2A lends support to the theory of a density gradient in the EW direction; its increased enhancement in the red lobe reflects the increased strength of the shocks in that lobe. \item \HCOplus\ is the least enhanced molecule and is generally confined to the base of the outflow, and the central driving source. This illustrates the effects of shocks on the chemistry in outflows (as traced by HCN), compared with that in protostellar cores and their envelopes (as traced by \HCOplus). \end{enumerate} A natural extension to this work would be to investigate all the other outflows in Perseus that are traced by both HCN and \HCOplus, to determine how typical IRAS~2A is as an outflow source, and whether other outflows in the region exhibit similar molecular enhancements. This is intended as the subject of a future paper. In conclusion, we find that disentangling the effects of the physical environments (density, temperature) of outflows, from the actual chemical processes occurring within the outflows is very difficult. A next step in the analysis would be the modelling of these protostellar sources and their outflows using 3D radiative transfer code, such as ARTIST \citep{2011IAUS..270..451P}, to better constrain the physical conditions giving rise to the emission we observe.	14	3	1403.3051	We present the results of a large-scale survey of the very dense (n > 10<SUP>6</SUP> cm<SUP>-3</SUP>) gas in the Perseus molecular cloud using HCO<SUP>+</SUP> and HCN (J = 4 → 3) transitions. We have used this emission to trace the structure and kinematics of gas found in pre- and protostellar cores, as well as in outflows. We compare the HCO<SUP>+</SUP>/HCN data, highlighting regions where there is a marked discrepancy in the spectra of the two emission lines. We use the HCO<SUP>+</SUP> to identify positively protostellar outflows and their driving sources, and present a statistical analysis of the outflow properties that we derive from this tracer. We find that the relations we calculate between the HCO<SUP>+</SUP> outflow driving force and the M<SUB>env</SUB> and L<SUB>bol</SUB> of the driving source are comparable to those obtained from similar outflow analyses using <SUP>12</SUP>CO, indicating that the two molecules give reliable estimates of outflow properties. We also compare the HCO<SUP>+</SUP> and the HCN in the outflows, and find that the HCN traces only the most energetic outflows, the majority of which are driven by young Class 0 sources. We analyse the abundances of HCN and HCO<SUP>+</SUP> in the particular case of the IRAS 2A outflows, and find that the HCN is much more enhanced than the HCO<SUP>+</SUP> in the outflow lobes. We suggest that this is indicative of shock enhancement of HCN along the length of the outflow; this process is not so evident for HCO<SUP>+</SUP>, which is largely confined to the outflow base.	false	[ "outflow properties", "similar outflow analyses", "protostellar outflows", "outflows", "driving force", "HCO", "the outflow properties", "the IRAS 2A outflows", "<", "HCN", "the outflow base", "the outflow lobes", "reliable estimates", "the outflow", "the outflows", "gas", "the driving source", "their driving sources", "young Class 0 sources", "SUP>12</SUP" ]	11.314852	10.765327	183
437554	[ "Koch, Andreas", "McWilliam, Andrew" ]	2014A&A...565A..23K	[ "The chemical composition of a regular halo globular cluster: NGC 5897" ]	31	[ "Zentrum für Astronomie der Universität Heidelberg, Landessternwarte, Königstuhl 12, 69117, Heidelberg, Germany ,", "Carnegie Observatories, 813 Santa Barbara St., Pasadena, CA, 91101, USA" ]	[ "2015A&A...578A.116C", "2015A&A...579A.104S", "2016A&A...587A.124K", "2016MNRAS.455.2417R", "2017A&A...599A..97H", "2017A&A...601A..41K", "2017A&A...607A..44B", "2018A&A...609A..13K", "2018A&A...619A.134H", "2018ApJ...862...71S", "2018MNRAS.478.1520B", "2019A&A...632A..55K", "2019A&ARv..27....8G", "2019AJ....158...14N", "2019ApJ...870...83J", "2019ApJ...883...84R", "2019MNRAS.483.1674R", "2020A&A...640A..87M", "2020ApJ...899...41S", "2021A&A...645A..64K", "2021A&A...646A...9C", "2021A&A...653A...2K", "2021A&A...655A..26L", "2021MNRAS.505.5957B", "2022Ap&SS.367...32S", "2022ApJ...934..150L", "2022AstL...48..578G", "2023A&A...677A..61M", "2023ApJ...944...58L", "2023AstL...49..673G", "2023MNRAS.525.3486J" ]	[ "astronomy" ]	6	[ "stars: abundances", "Galaxy: abundances", "Galaxy: evolution", "globular clusters: individual: NGC 5897", "Galaxy: halo", "Astrophysics - Astrophysics of Galaxies" ]	[ "1968ApJ...153..569S", "1973PhDT.......180S", "1976ApJS...30....1V", "1980A&A....92...70R", "1981A&A....97..391T", "1982A&A...107....1N", "1984ApJS...55...45Z", "1988AJ.....96...92A", "1990A&A...228..155N", "1990AJ.....99..250W", "1992AJ....104..627G", "1993AJ....105.1145F", "1993ApJ...419..207R", "1995AJ....109..605G", "1995AJ....109.2757M", "1995AcA....45..653J", "1995ApJS...98..617T", "1997A&AS..121...95C", "1997MNRAS.287..481W", "1997PASP..109..883R", "1997PASP..109..907R", "1998AJ....115.1640M", "1998ApJ...500..525S", "2000AJ....120.1351S", "2000ApJ...531..159K", "2000ApJ...544..302B", "2001AJ....121..916T", "2001AJ....122.1464C", "2001AJ....122.2587C", "2001ApJ...557..802B", "2003A&A...402..985Y", "2003AJ....125..684S", "2003PASP..115..143K", "2003PASP..115..413S", "2003PASP..115..688K", "2004AJ....128.1177V", "2004ApJ...601..864T", "2004ApJ...612..168P", "2004ApJ...617.1091S", "2004PASJ...56.1041S", "2005A&A...438..875Y", "2005AJ....130.2140P", "2005ApJ...626..465R", "2007A&A...464.1029D", "2007ApJ...661.1152F", "2008AJ....135.1551K", "2008ApJS..178...89D", "2008MNRAS.391..825D", "2009A&A...500.1143F", "2009A&A...503..545L", "2009A&A...505..117C", "2009A&A...505..139C", "2009A&A...506..729K", "2009A&A...508..695C", "2009ARA&A..47..371T", "2009ARA&A..47..481A", "2009ApJ...696..797C", "2009ApJ...696..870D", "2009IAUS..258..233P", "2009PASJ...61..563T", "2010A&A...513A..34S", "2010A&A...516A..55C", "2010A&A...517A..59K", "2010A&A...517A..81G", "2011A&A...527A..65H", "2011A&A...528A.103L", "2011A&A...530A..15N", "2011AJ....141..175S", "2011AJ....142...63K", "2012A&ARv..20...50G", "2012AJ....143...48H", "2012MNRAS.426.2889M", "2012MNRAS.427...27B", "2012MNRAS.427...50L", "2013A&A...554A..81K", "2013A&A...557A..26M", "2013A&A...557A.128C", "2013A&A...560A...8M", "2013ApJ...771...67I", "2013ApJS..205...11L", "2013ApJS..208...27W" ]	[ "10.1051/0004-6361/201323119", "10.48550/arXiv.1403.1262" ]	1403	1403.1262_arXiv.txt	Globular clusters (GCs) are amongst the oldest stellar systems in the Universe and purvey important information on the earliest evolutionary stages of the Galaxy. As such, they have received much attention over the past years as they also harbor a multitude of fascinating properties. Long believed to be {\em the} classical example of simple stellar populations, it is now accepted that GCs consist of multiple stellar populations that are evident in their color-magnitude diagrams (CMDs; e.g., Piotto 2009, and references therein) and their light chemical element distributions (e.g., Gratton et al. 2012, and references therein. These are a direct consequence of the presence of at least two generations of stars, differing in age and chemical abundances, that formed within a few 100 Myr of each other. Stars of the first generations (either fast rotating massive stars, Decressin et al. 2007, or intermediate-mass Asymptotic Giant Branch [AGB] stars, e.g., D'Ercole et al. 2008) pollute the Interstellar Medium (ISM) with the products of p-capture reactions, out of which the dominant, second stellar generation is born. This enrichment scenario leaves the $\alpha$- and Fe-peak elements unaltered (as is indeed observed in GCs), while producing the characteristic light-element variations such as the Na-O and Mg-Al anti-correlations and CN and CH bimodalities. Considering the broad variety in the GCs' properties (horizontal branch [HB] morphology, central concentration, environment, to name a few), it is then striking that several key features are found throughout every single GC studied to date. For instance, several authors have now resorted to {\em defining} a GC as an object that shows a Na-O correlation (Carretta et al. 2010). In this work, we will add the poorly-studied inner halo GC NGC~5897 to the family of typical GCs by measuring in-depth for the first time its chemical composition. This will reveal if also NGC~5897 exhibits the canonical chemical abundance variations and how these could relate to the factors of, e.g., environment. NGC~5897 is a metal-poor GC in the Galactic halo, located at a moderate distance from the Sun (R$_{\odot}$=12.5 kpc; R$_{\rm GC}$=7 kpc; Harris 2006 [2010 edition].). It shows a predominantly blue HB with a morphology that has been labeled ``normal'' considering its low metallicity (e.g., Clement \& Rowe 2001; Fig.~1). However, some of its other properties still defy a satisfying explanation: In a photometric study, Testa et al. (2001) analysed NGC~5897 and three other GCs, which were all confirmed to be coeval and to have similar metallicity. Thus it was concluded that there must exist a very prominent, {additional} parameter besides age and metallicity governing the differences in their HB morphologies. This parameter was taken to be ``environment'', represented, e.g., by the concentration or central densities of such stellar systems ({Fusi-Pecci et al. 1993}). In this context, NGC~5897 has a very low concentration (half-light and tidal radii are 2.1$\arcmin$ and 10.1$\arcmin$, respectively); its concentration parameter (c=$\log\,[r_t/r_c]$) is 0.86 and, in fact, only 15\% of the entire known Milky Way GCs are less densely concentrated, where 60\% of those are within the inner halo at R$_{\rm GC}<20$ kpc (Harris 1996). {Further, possible "third parameters" comprise the Helium abundance of the clusters (Gratton et al. 2010).} It is very surprising that, given the wealth of photometric data for this GC, no consensus as to its metallicity is in sight, yet: the current values for [Fe/H] in the literature, obtained from photometry or low-resolution spectroscopy, range from $-1.68$ (Zinn \& West 1984) to $-2.09$ dex (Kraft \& Ivans 2003), which is by far a larger uncertainty than the possible accuracy achievable with present-day spectroscopic methods. Such an ignorance would also seem to prohibit an accurate assessment even of the simplest parameters governing the HB morphology. NGC~5897 is known to harbour several variable stars, mainly RR~Lyrae. Their periods were measured to be longer than 0.6 days, which is not unusual for a low-metallicity cluster. Puzzling is, however, the unparalleled, long mean period of the RR$ab$ variables of $\left<P_{ab}\right> >$ 0.8 d (e.g., Wehlau 1990; Clement et al. 2001; Clement \& Rowe 2001), which could indicate even lower metallicities than the mean of this GC. This work is organized as follows: In \S2 we describe our target selection and data gathering; the abundance and error analyses are laid out in detail in \S\S3,4. All our results are presented and interpreted in \S5 before summarizing our findings in \S6.	Our abundance analysis of a large number of chemical tracers in the metal-poor inner halo cluster NGC~5897 identified this GC as fully representative of the bulk of such objects in the Milky Way. This includes the enhancement in the $\alpha$-elements to $\sim$0.3 dex, approximately Solar Fe-peak elements, and n-capture elements that were predominantly produced by the $r$-process. The later notion is consistent with only negligible AGB contributions to the chemical inventory in the metal poor (read: early-time) regime. The presence of a pronounced Na-O anti-correlation renders this object, by definition, a {\em globular} cluster (not that there ever was a doubt). Unfortunately, our limited sample size did not allow us to trace the full extent of this relation, in particular we are lacking information about the spatial variations of the first vs. second generations, but the severe crowding prohibited exposures of the central stars. All other elements, safe for O, Na, and Al only show little star-to-star scatter, mostly explicable by the measurement uncertainties or in part driven by peculiar patterns in individual stars. Several features such as the shape of the Na-O correlation and the preponderance of the $r$-process is reminiscent of the metal poor ([Fe/H]$\sim -2.6$ dex) halo GC M15. Our work has revealed NGC~5897 to be more metal poor than implied by older CMD and low-resolution metallicity studies. We suggest that this lower metallicity can explain the remarkably longer periods of the RR Lyr found in this GC. Indeed, the error-weighted mean metallicity of the RR Lyr of Clement \& Rowe (2001) is [Fe/H] = $-2.2\pm0.2$ on the scale of Zinn \& West (1984), where we adopted the relations between period, phase, and metallicity of Jurcsik (1995); see also Haschke et al. (2012). While the errors on individual stars are tremendous, the mean value of the RR\,Lyr is thus in line with the metal poor nature of this GC found from our red giants.	14	3	1403.1262	We report for the first time on the chemical composition of the halo cluster NGC 5897 (R<SUB>⊙</SUB> = 12.5 kpc), based on chemical abundance ratios for 27 α-, iron-peak, and neutron-capture elements in seven red giants. From our high-resolution, high signal-to-noise spectra obtained with the Magellan/MIKE spectrograph, we find a mean iron abundance from the neutral species of [Fe/H] = - 2.04 ± 0.01 (stat.) ± 0.15 (sys.), which is more metal-poor than implied by previous photometric and low-resolution spectroscopic studies. The cluster NGC 5897 is α-enhanced (to 0.34 ± 0.01 dex) and shows Fe-peak element ratios typical of other (metal-poor) halo globular clusters (GCs) with no overall, significant abundance spreads in iron or in any other heavy element. Like other GCs, NGC 5897 shows a clear Na-O anti-correlation, where we find a prominent primordial population of stars with enhanced O abundances and approximately solar Na/Fe ratios, while two stars are Na-rich, providing chemical proof of the presence of multiple populations in this cluster. Comparison of the heavy element abundances with the solar-scaled values and the metal-poor GC M15 from the literature confirms that NGC 5897 has experienced little contribution from s-process nucleosynthesis. One star of the first generation stands out in that it shows very low La and Eu abundances. Overall, NGC 5897 is a well behaved GC showing archetypical correlations and element-patterns, with little room for surprises in our data. We suggest that its lower metallicity could explain the unusually long periods of RR Lyr that were found in NGC 5897. <P />This paper includes data gathered with the 6.5-m Magellan Telescopes located at Las Campanas Observatory, Chile.Table 5 is available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201323119/olm">http://www.aanda.org</A>Full Table 2 is only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A23">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A23</A>	false	[ "chemical abundance ratios", "enhanced O abundances", "significant abundance", "s-process nucleosynthesis", "halo globular clusters", "NGC", "Las Campanas Observatory", "Magellan Telescopes", "iron", "other GCs", "multiple populations", "the heavy element abundances", "a mean iron abundance", "little room", "Chile", "little contribution", "Fe-peak element ratios", "data", "stars", "Magellan" ]	8.913518	8.746373	140
814096	[ "Dvorkin, Cora", "Wyman, Mark", "Rudd, Douglas H.", "Hu, Wayne" ]	2014PhRvD..90h3503D	[ "Neutrinos help reconcile Planck measurements with both the early and local Universe" ]	122	[ "Institute for Advanced Study, School of Natural Sciences, Einstein Drive, Princeton, New Jersey 08540, USA", "New York University, Center for Cosmology and Particle Physics, New York, New York 10003, USA", "Kavli Institute for Cosmological Physics, Research Computing Center, University of Chicago, Chicago, Illinois 60637, USA", "Kavli Institute for Cosmological Physics, Department of Astronomy and Astrophysics, Enrico Fermi Institute, University of Chicago, Chicago, Illinois 60637, USA" ]	[ "2014ApJ...794..135B", "2014EPJC...74.2954Z", "2014JCAP...06..031A", "2014JCAP...06..042A", "2014JCAP...06..061H", "2014JCAP...07..006G", "2014JCAP...08..036M", "2014JCAP...08..038L", "2014JCAP...08..048H", "2014JCAP...09..035N", "2014JCAP...10..025H", "2014JCAP...10..044Z", "2014JCAP...10..053B", "2014JCAP...11..039B", "2014JHEP...10..104B", "2014JHEP...10..184A", "2014JHEP...11..018D", "2014JHEP...12..104K", "2014MNRAS.442..670C", "2014MNRAS.445..778I", "2014PhLB..738..206L", "2014PhLB..738..254F", "2014PhLB..739...62K", "2014PhLB..739..102Z", "2014PhRvD..89l1302C", "2014PhRvD..90b3543D", "2014PhRvD..90b3544H", "2014PhRvD..90d3001S", "2014PhRvD..90d3534D", "2014PhRvD..90f3007B", "2014PhRvD..90g3002A", "2014PhRvD..90j3512B", "2014PhRvD..90l3524H", "2014PhRvL.113c1301S", "2014PhRvL.113d1301L", "2014PhRvL.113p1301C", "2014SCPMA..57.1431L", "2014SCPMA..57.1455L", "2014arXiv1404.3804X", "2014arXiv1404.6160G", "2014arXiv1405.2775Z", "2014arXiv1407.8105A", "2015AJ....149..183H", "2015ApJ...812...25M", "2015IJMPD..2441001C", "2015JCAP...01..017M", "2015JCAP...01..034T", "2015JCAP...01..040K", "2015JCAP...03..046F", "2015JCAP...04..038Z", "2015JCAP...07..014F", "2015JPhCS.598a2010V", "2015MNRAS.446.2205M", "2015MNRAS.446.2871C", "2015MNRAS.448.3463C", "2015MNRAS.451.2877M", "2015PhLB..740..359Z", "2015PhLB..744..213L", "2015PhLB..750..201T", "2015PhRvD..91e5033A", "2015PhRvD..91f3009R", "2015PhRvD..91f3504M", "2015PhRvD..91f5021A", "2015PhRvD..91i5023B", "2015PhRvD..91l3505D", "2015PhRvD..92f3501M", "2015PhRvD..92l5027G", "2015arXiv150607433R", "2015ist..book.....B", "2016ARNPS..66..401A", "2016ApJ...826...56R", "2016EPJC...76..588X", "2016JCAP...02..059B", "2016JCAP...10..019B", "2016PhRvD..93h3011Z", "2016PhRvD..94h3519W", "2016arXiv160309102G", "2017EPJC...77..418F", "2017EPJC...77..725W", "2017EPJWC.15207001M", "2017FrP.....5...53A", "2017MNRAS.471.2254Z", "2017PhLB..764..322B", "2017PhR...711....1A", "2017PhRvD..95c5028G", "2017PhRvD..96d3520Z", "2017PhRvD..96h3526O", "2017SSRv..212.1817S", "2017arXiv170805372C", "2018JPhG...45i5003S", "2018PASP..130h4502Z", "2018PhLB..779..473Z", "2018PhRvD..98h3525P", "2018SCPMA..61e0411F", "2018SSRv..214...90K", "2019A&A...631A..96I", "2019EPJC...79..557R", "2019MNRAS.484..409R", "2019PDU....23..261F", "2019PhRvD.100f3542L", "2019PhRvD.100j3513A", "2019adds.book...35S", "2019adds.book..219K", "2020MNRAS.492.5052Y", "2020PhRvD.101d3518Z", "2020PhRvD.102f3527N", "2020PhRvD.102l3523L", "2020SCPMA..6320401F", "2020SCPMA..6390404F", "2021JCAP...04..024Y", "2021PDU....3100762Y", "2021PhRvD.103b1301W", "2021PhRvD.103l1301C", "2022A&A...666A..34S", "2022EPJC...82...78Y", "2022JCAP...04..048S", "2022NewAR..9501659P", "2022PhRvD.106j3517E", "2022Univ....8..284S", "2022arXiv220814435G", "2023PhRvD.107f3536Z", "2024arXiv240510358B" ]	[ "astronomy", "physics" ]	6	[ "98.80.-k", "95.85.Ry", "98.70.Vc", "98.80.Es", "Cosmology", "Neutrino muon pion and other elementary particles", "cosmic rays", "Background radiations", "Observational cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Phenomenology" ]	[ "1993PhRvD..47.5219B", "2002PhRvD..66j3511L", "2006PhRvD..74l7305B", "2010PhRvD..81f7302M", "2011ApJ...730..119R", "2011ApJ...743...28K", "2011MNRAS.418.1707B", "2012ApJ...755...70R", "2012AstL...38..347B", "2012JHEP...12..123G", "2012MNRAS.427.2132P", "2012arXiv1204.5379A", "2013A&A...558A..57I", "2013ARNPS..63...45C", "2013ApJS..208...20B", "2013JCAP...09..013V", "2013JCAP...10..044H", "2013MNRAS.428.1036M", "2013arXiv1302.5086R", "2014A&A...571A..16P", "2014ApJ...782...74H", "2014JCAP...03..011V", "2014JCAP...04..014D", "2014JCAP...05..046K", "2014JCAP...07..014C", "2014JCAP...08..039F", "2014JCAP...08..053A", "2014JCAP...10..035M", "2014MNRAS.438...49R", "2014MNRAS.440.1138E", "2014PhLB..734..167C", "2014PhLB..736..305D", "2014PhRvD..89j1302M", "2014PhRvD..90d3507G", "2014PhRvL.112e1302W", "2014PhRvL.112e1303B", "2014PhRvL.112q1301L", "2014PhRvL.112q1302M", "2014PhRvL.112s1303B", "2014PhRvL.112x1101B", "2014PhRvL.113d1301L", "2015PhLB..740..359Z", "2015PhRvD..91b3518S" ]	[ "10.1103/PhysRevD.90.083503", "10.48550/arXiv.1403.8049" ]	1403	1403.8049_arXiv.txt	\label{sec:modelsdata} The simplest inflationary $\Lambda$CDM model is characterized by 6 parameters, $\{\Oc h^2, \, \Ob h^2, \, \tau, \theta_{\rm MC}, A_{s}, n_s \}$. Here, $\Oc h^2$ represents the physical cold dark matter (CDM) density, $\Ob h^2$ is the baryon density, $\tau$ is the Thomson optical depth to reionization, $\theta_{\rm MC}$ is a proxy for the angular acoustic scale at recombination, $A_s$ is the amplitude of the initial curvature power spectrum at $k = 0.05$ Mpc$^{-1}$, and $n_s$ its spectral index. To these parameters we add $r$, the tensor-scalar ratio evaluated at $k=0.002$ Mpc$^{-1}$, and take the tensor tilt to follow the consistency relation $n_t = -r/8$. For the neutrino extension we add two new parameters. The first is the effective number of relativistic species, $N_{\rm eff}$, defined in terms of the relativistic energy density at high redshift \be \rho_{\rm r} = \rho_\gamma + \rho_\nu = \left[ 1 + \frac{7}{8} $\frac{4}{11}$^{4/3} \hspace{-6pt} N_{\rm eff} \right] \rho_\gamma. \ee In the minimal model, $N_{\rm eff}=3.046$ and so $\Delta N_{\rm eff}=N_{\rm eff}-3.046>0$ indicates the presence of extra relativistic particle species in the early Universe. We assume that the active neutrinos have $\sum m_\nu=0.06$ eV (as suggested by a normal hierarchy and solar and atmospheric oscillation measurements \cite{GonzalezGarcia:2012sz}) and any additional contributions are carried by a mostly sterile state, with effective mass $m_s$ for a total neutrino contribution to the energy density today of \be \label{omegan} ({94.1 \, \rm eV}) \Omega_\nu h^2 = {(3.046/3)^{3/4} \sum m_\nu+m_s}. \ee Thus $m_s$ characterizes extra non-CDM energy density rather than the true (Lagrangian) mass of a neutrino-like particle. In particular even in the \dneff$\rightarrow 0$ limit, at fixed $m_s$, the presence of an extra sterile species still adds an extra energy density component at recombination and still changes inferences based on the sound horizon. In the most recent version of CosmoMC \cite{Lewis:2002ah}, a prior on the physical mass of a thermally produced sterile neutrino is imposed with a value of $m_s/(\Delta N_{\rm eff})^{3/4} < 10{\rm eV}$ to close off a degeneracy between very massive neutrinos and cold dark matter. We call these extensions the \LCDMn\ and \LCDMrn\ models, where the names follow from the additional parameters introduced by tensors and neutrinos. \begin{table}[t] \centering \begin{tabular}{ c \| c } & Models\\ \hline \LCDM & $\{\Oc h^2, \, \Ob h^2, \, \tau, \theta_{\rm MC}, A_{s}, n_s \}$ \\ \LCDMr & \LCDM + r \\ \LCDMn & \LCDM +\neff$+m_s$ \\ \LCDMrn & \LCDM +\neff$+m_s+r$\\ \multicolumn{2}{c}{} \\ & Data sets\\ \hline \M & \{Planck, WP, SPT/ACT\} \\ \EU & \M +BICEP2 \\ \LU & \M +$H_0$+BAO+Clusters\\\ \ELU &\M +BICEP2 +$H_0$+BAO+Clusters\\ \end{tabular} \caption {Model and data set combinations. All models include \LCDM\ parameters. All data sets include the \midd\ Universe (\M) set. } \label{tab:datamodel} \end{table} For the data sets, we consider combinations that best expose the separate early and local Universe tensions with what we refer to as the \midd\ (\M) Universe -- the time intermediate between inflation at the late Universe, characterized chiefly by the CMB temperature power spectrum. This \M\ data set is composed of the Planck temperature \cite{Ade:2013zuv}, WMAP9 polarization \cite{Bennett:2012fp}, and ACT/SPT \cite{Das:2013zf,Reichardt:2011yv,Keisler:2011aw} high multipole power spectra. For the Planck data analysis we marginalize over the standard Planck foreground parameters \cite{Ade:2013zuv}. For the early Universe tension, we add to these what we call the early (E) data set, the BICEP2 BB and EE polarization bandpowers \cite{Ade:2014xna}. We call the combination of these two the early-\midd\ (\EU) Universe data set. For the local Universe tension, we define the following collection as the local (L) data sets: the $H_0$ inference from the maser-cepheid-supernovae distance ladder, $h = 0.738 \pm 0.024$ \cite{Riess:2011yx}, BAO measurements \cite{Anderson:2012sa, Padmanabhan:2012hf, Blake:2011en}, and the X-ray derived cluster abundance using the likelihood code\footnote{This prescription employs the total matter power spectrum. More recent studies of the cluster abundance indicate that this may somewhat underestimate the neutrino mass at least at lower masses than considered here \cite{Villaescusa-Navarro:2013pva} where the abundance follows the cold dark matter power spectrum more closely. We continue to adopt this approach to be conservative with respect to new neutrino physics and provide a continuous limit with CDM at high $m_s$.} of Ref.~\cite{Burenin:2012uy} which roughly equates to a constraint on $S_8=\sigma_8 (\Omega_m/0.25)^{0.47}=0.813 \pm 0.013$. Note that the BAO data are added here not because they are in tension with Planck (they are not) but because they exclude resolutions of the $H_0$ tension involving exotic dark energy or curvature. To these we again add the Planck temperature, WMAP9 polarization and the ACT/SPT data sets. We call this the \midd-local (\LU) Universe data set. Finally, we call the union of this with the \EU\ data the \ELU\ data set. These model and data choices are summarized in Tab.~\ref{tab:datamodel}. We analyze these data and models using the Markov Chain Monte Carlo technique and the CosmoMC code \cite{Lewis:2002ah}. Our local analysis will be in many ways similar to that performed in Ref.~\cite{Wyman:2013lza} to which we refer the reader for details and robustness checks. \begin{table}[t] \centering \begin{tabular}{ c \| c \| c\| c } & \LCDMrn--\EU & \LCDMn--\LU & \LCDMrn--\ELU \\ \hline \dneff\ &$1.06\pm0.37$ & $0.62\pm0.28$ & $0.98\pm0.26$ \\ $m_s$ [eV] &$<0.22$ & $0.48\pm0.15$ & $0.52\pm0.13$ \\ $r$ & $0.19\pm0.05$ & -- & $0.22\pm0.05$ \\ \hline $100\Omega_b h^2$&$2.268\pm0.043$ & $2.267\pm0.028$ & $2.276\pm0.027$\\ $\Omega_c h^2$ &$0.132\pm0.005$ & $0.122\pm0.005$ & $0.127\pm0.004$\\ $100\theta_{\rm MC}$&$1.040\pm0.001$ & $1.041\pm0.001$ & $1.040\pm0.001$\\ $\tau$ &$0.100\pm 0.015$ & $0.096\pm0.014$ & $0.097\pm0.014$\\ $\ln (10^{10} A_s)$ & $3.136\pm0.033$ & $3.107\pm0.031$ & $3.117\pm0.030$\\ $n_s$ &$0.999\pm0.017$ & $0.985\pm0.012$ & $1.001\pm0.010$\\ \hline $h$ & $0.74\pm0.04$ & $0.70\pm0.01$ & $0.72\pm0.01$\\ $S_8$ &$0.89\pm0.03$ & $0.81\pm0.01$ & $0.81\pm0.01$\\ \end{tabular} \caption {Parameter constraints (68\% confidence level) with various model and data assumptions. Note that the \LCDMn-\LU\ case is in a different, no tensor model, context than the others which affects parameter interpretations.} \label{tab:results} \end{table}	\label{sec:discussion} We have shown that the tension introduced by the detection of large amplitude gravitational wave power by the BICEP2 experiment with temperature anisotropy measurements is alleviated in a model with an extra sterile neutrino. The relativistic energy density required to alleviate this early Universe tension is the same as that required to resolve the late Universe tension of acoustic distance measures with local Hubble constant measurements. Combined they imply a \dneff$=0.98\pm 0.26$. Note that the ACT/SPT data already limit the upper range of allowed \dneff, and so this explanation of the early and late Universe tensions can be tested with more data from high multipole CMB temperature and polarization observations. Conversely, it would weaken if extra evidence for alternate inflationary or foreground explanations of the early Universe tension is found. By making the sterile neutrino massive, the tension with growth of structure measurements can simultaneously be alleviated. The combined constraint on the effective mass is $m_s=0.52\pm 0.13$ eV. This preference for high neutrino mass(es) is mainly driven by the cluster data set (cf.~Ref.\ \cite{Verde:2013cqa} who find upper limits without clusters). Compared with analyses that do not include gravitational waves, the sterile neutrino suggested here has a smaller expected physical mass, thanks to the generally larger values of \dneff\ that we find (recall the thermal conversion formula: $m_s^{\rm th} = m_s/(\Delta N_{\rm eff})^{3/4}$). At the upper range of \dneff, it is thus somewhat less likely that this sterile neutrino could explain anomalies in short baseline and reactor neutrino experiments (see Refs.\ \cite{Conrad:2013mka,Abazajian:2012ys} for reviews). If future data or analyses lead to increased mass estimates for the clusters, that change would weaken the preference for a non-zero $m_s$ by increasing $\som$, giving better concordance with the basic \LCDM\ prediction. However, the preference can only be eliminated if the systematic shift is roughly triple the 9\% estimate; that estimate is derived from comparisons of a variety of X-ray, optical, Sunyaev-Zel'dovich, and lensing observables (see e.g.~\cite{Rozo:2013hha} for a recent assessment). Taken at face value, these results leave us with a potentially very different cosmological standard model. Gravitational waves are nearly indisputable evidence for an inflationary epoch, but the lack of a significant primordial tilt compared with $r$ would suggest somewhat unusual inflationary physics is at play where the impact of the slope and curvature of the inflationary potential partially cancelled each other. On the other hand, the new neutrino physics favored here is less contrived than that proposed in Ref.~\cite{Wyman:2013lza,Hamann:2013iba,Battye:2013xqa}: we are now allowed \dneff$=1$, which is in better accord with a ``theory prior" that the sterile neutrino would be fully populated by oscillations with active neutrinos for typical mixing angles. Meanwhile, we see somewhat less residual tension between the early and late Universe, especially if galaxy clusters are indeed a bit more massive than we have assumed in our main analysis. Each of these new ingredients will soon be cross checked by a wide variety of upcoming observations. If all are confirmed, observational cosmology will have provided not one but two clear discoveries of particle physics beyond the Standard Model within a short space of time, giving long sought clear guidance for how to advance physics into the future.	14	3	1403.8049	In light of the recent BICEP2 B-mode polarization detection, which implies a large inflationary tensor-to-scalar ratio r<SUB>0.05</SUB>=0.2-0.05+0.07, we re-examine the evidence for an extra sterile massive neutrino, originally invoked to account for the tension between the cosmic microwave background (CMB) temperature power spectrum and local measurements of the expansion rate H<SUB>0</SUB> and cosmological structure. With only the standard active neutrinos and power-law scalar spectra, this detection is in tension with the upper limit of r<0.11 (95% confidence) from the lack of a corresponding low multipole excess in the temperature anisotropy from gravitational waves. An extra sterile species with the same energy density as is needed to reconcile the CMB data with H<SUB>0</SUB> measurements can also alleviate this new tension. By combining data from the Planck and ACT/SPT temperature spectra, WMAP9 polarization, H<SUB>0</SUB>, baryon acoustic oscillation and local cluster abundance measurements with BICEP2 data, we find the joint evidence for a sterile massive neutrino increases to ΔN<SUB>eff</SUB>=0.98±0.26 for the effective number and ms=0.52±0.13 eV for the effective mass, or 3.8σ and 4σ evidence, respectively. We caution the reader that these results correspond to a joint statistical evidence and, in addition, astrophysical systematic errors in the clusters and H<SUB>0</SUB> measurements, and small-scale CMB data could weaken our conclusions.	false	[ "H<SUB>0</SUB", "local cluster abundance measurements", "local measurements", "evidence", "BICEP2 data", "gravitational waves", "tension", "CMB", "data", "baryon acoustic oscillation", "cosmological structure", "WMAP9 polarization", "small-scale CMB data", "a sterile massive neutrino increases", "a joint statistical evidence", "H<SUB>0</SUB> measurements", "the CMB data", "the expansion rate H<SUB>0</SUB", "the joint evidence", "BICEP2" ]	7.068111	-1.398102	41
437500	[ "Moriya, Takashi J." ]	2014A&A...564A..83M	[ "Mass loss of massive stars near the Eddington luminosity by core neutrino emission shortly before their explosion" ]	32	[ "Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, 53121, Bonn, Germany ; Research Center for the Early Universe, Graduate School of Science, University of Tokyo, Hongo 7-3-1, Bunkyo, 113-033, Tokyo, Japan" ]	[ "2014ApJ...796..121J", "2014ApJ...797..118F", "2014MNRAS.445.2492M", "2014Natur.512..282M", "2015ApJ...803L..26M", "2016MNRAS.455..423M", "2016MNRAS.459L..21K", "2016PhDT.......149T", "2017A&A...599A.129T", "2017ApJ...851..138D", "2017IAUS..331...11D", "2017MNRAS.464.3249S", "2017MNRAS.469L.108M", "2017MNRAS.470.1642F", "2017hsn..book..843B", "2018MNRAS.476.2840M", "2019ApJ...884...87T", "2019MNRAS.482.2277D", "2020A&A...635A..39T", "2020A&A...635A.127K", "2020ApJ...897L..44T", "2021ApJ...907...99S", "2021MNRAS.503.5965Y", "2022ApJ...929..177T", "2022MNRAS.515.3597L", "2023ApJ...952L..30J", "2023ApJ...953L..16H", "2024ApJ...961...47I", "2024ApJ...966...30T", "2024MNRAS.528.4209M", "2024arXiv240112403A", "2024arXiv240503747R" ]	[ "astronomy" ]	5	[ "stars: massive", "stars: mass-loss", "supernovae: general", "stars: evolution", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1980Ap&SS..69..115N", "1988A&AS...72..259D", "1996ApJS..102..411I", "2001A&A...369..574V", "2001A&A...373..555M", "2004APh....21..303O", "2006ApJ...644.1151S", "2007Natur.447..829P", "2007Natur.450..390W", "2009Natur.458..865G", "2011ApJ...727..104C", "2011ApJS..192....3P", "2011MNRAS.414.1715B", "2011MNRAS.414.2985D", "2011MNRAS.415..773S", "2011Natur.474..487Q", "2011NuPhS.221..380O", "2012A&A...541A.129L", "2012ARA&A..50..107L", "2012ApJ...744...10K", "2012ApJ...756L..22M", "2012ApJ...760..154C", "2012MNRAS.423L..92Q", "2013A&A...555A..10T", "2013ApJ...763L..27P", "2013ApJ...767....1P", "2013ApJ...767..143B", "2013ApJ...769..109L", "2013ApJ...773...76C", "2013ApJ...778L..23T", "2013ApJ...779L...8F", "2013ApJS..208....4P", "2013MNRAS.430.1801M", "2013MNRAS.433.1312F", "2013MNRAS.435.1520M", "2013Natur.494...65O", "2014ApJ...780...21M", "2014ApJ...780...96S", "2014ApJ...781...13L", "2014ApJ...785...37B", "2014MNRAS.438.1191S", "2014MNRAS.441..289B", "2014MNRAS.443..671G" ]	[ "10.1051/0004-6361/201322992", "10.48550/arXiv.1403.2731" ]	1403	1403.2731_arXiv.txt	Stellar mass loss is one of the most influential phenomena which determine the fate of a massive star. For instance, the spectral type of an exploding star, a supernova (SN), depends on the amount of mass loss during the evolution. There is a growing evidence that some stars experience very high mass-loss rates shortly before the explosion. SN 2006jc is, for the first time, found to have got bright two years before its explosion as a Type Ibn SN \citep{pastorello2007}. A similar Type Ibn SN iPTF13beo is found to have two light curve peaks which indicate that the mass-loss rate of the Wolf-Rayet (WR) progenitor is $\sim 2\times 10^{-3}$ $M_\odot~\mathrm{yr^{-1}}$ in a year before the explosion \citep{gorbikov2013}. Type IIn SNe are more commonly discovered SNe with dense circumstellar media (CSM) resulting from the high mass-loss rates of their progenitors shortly before their explosions. They are typically higher than $\sim 10^{-3}$ $M_\odot~\mathrm{yr^{-1}}$ \citep[e.g.,][]{taddia2013,kiewe2012,moriya2013}. Having such high mass-loss rates immediately before the explosions of these SN progenitors is not expected from the standard stellar evolution theory and its mechanisms are still not known well. Some of them are likely to be related to the progenitors' very high luminosities and their proximity to the Eddington luminosity. The progenitor of Type IIn SN 2005gl is found to have an absolute brightness of $M_V\simeq-10$ mag and it belongs to the most luminous class of stars. The progenitor of Type IIn SN 2009ip provided us with more remarkable information about the progenitors' activities (e.g., \citealt{margutti2013,smith2013,mauerhan2013,prieto2013}, but see also \citealt{fraser2013,pastorello2013}). The luminosity of the progenitor was, again, close to $M_V=-10$ mag and it is presumed to be near the Eddington luminosity. The progenitor started to show several outbursts from a few years before its explosion. The final burst in August 2012 started about 50 days before the major luminosity increase started in September 2012 which may be caused by the SN explosion. A dense CSM could have been formed during these eruptive events. Type IIn SNe 2010mc and 2011ht also experienced a similar pre-SN outburst \citep{ofek2013,fraser2013b}. A Type Ic superluminous SN 2006oz had a precursor from about 10 days before its major luminosity increase \citep{leloudas2012}. This precursor may also be originated from the progenitor activities shortly before the explosion (but see also \citealt{moriya2012} for another interpretation). In addition to the unexpected mass-loss enhancement shortly before the core collapse, some stars may have larger radii than those expected from the stellar evolution theory at the time of the explosion. The early emission of Type Ib SN 2008D is suggested to be difficult to be explained by standard WR stars \citep[e.g.,][]{dessart2011}. \citet{bersten2013} showed that the radius of the WR progenitor of SN 2008D should be $\sim 10~R_\odot$ to explain the early emission by the adiabatic cooling of the SN ejecta and they concluded that $^{56}$Ni may have pumped out to the outermost layers of the WR progenitor (see also, e.g., \citealt{couch2011,balberg2011} for other interpretations). Some fast evolving Type Ic SNe are suggested to originate from extended WR stars whose radii are $\sim 10~R_\odot$ (\citealt{kleiser2013}, but see also \citealt{tauris2013}). The progenitors of ultra-long gamma-ray bursts are unlikely to be WR stars with their typical radii because of their long durations \citep[e.g.,][]{levan2013}. Those observations may infer that some WR stars explode with much larger radii than those predicted by the current stellar evolution theory. One of the characteristics of massive stars shortly before the explosion is their huge neutrino luminosities. Neutrinos are constantly emitted from the stellar core throughout the evolution because of the nuclear reactions. However, after the onset of the carbon burning, the thermal neutrino emission becomes significant because of the high temperature required for the ignitions of heavy elements. The expected neutrino luminosity is so large that dying stars will be detected by the modern neutrino detectors from a few weeks before the core collapse if they are close enough to the Earth \citep{odrzywolek2011,odrzywolek2004}. The high neutrino luminosity from the core together with the energy release by nuclear burning and the composition gradients results in the active convective motion which is suggested to be related to the high mass-loss rates of SN progenitors \citep[e.g.,][]{quataert2012,shiode2013}. Here, in this paper, we suggest another possible effect caused by the huge neutrino luminosity which may result in the extreme mass loss at the stellar surface. The mass loss of the core due to the neutrino emission leads to the sudden decrease of the gravitational force. The effect of the mass loss due to large neutrino luminosities after the core collapse of a star has been discussed \citep[][]{lovegrove2013,nadezhin1980}. In this paper, we suggest that the effect of the neutrino mass loss at the stellar core can appear at the stellar surface in massive stars near the Eddington luminosity immediately before their core collapse. In short, the gravitational potential can be significantly reduced because of the neutrino emission shortly before the core collapse and the stellar luminosity can exceed the Eddington luminosity. Thus, the weakened gravitational force results in the enhanced mass loss. It may also result in the expansion of massive stars shortly before their explosion.	We suggest a new mechanism to induce high mass-loss rates shortly before the explosion of massive stars. The neutrino luminosities of massive stars increase as they advance the nuclear burning. Neutrinos escape freely from the stellar core and take mass out of the core. The core mass-loss rate by the neutrino emission becomes higher than $10^{-4}$ $M_\odot~\mathrm{yr^{-1}}$ at about a year before the core collapse. The core mass-loss rate becomes higher than $10^{-2}$ $M_\odot~\mathrm{yr^{-1}}$ in $\sim10$ days before the core collapse (Figure \ref{neulumi}). If a star is close enough to the Eddington luminosity at that time, the stellar luminosity can exceed the Eddington luminosity by the core mass loss because of the reduced gravitational force. Then, the radiative force which was balanced by the gravitational force can push the material at the stellar surface outward and drive strong mass loss. As the reduced gravitational energy is available to push the surface meterial, the mass-loss rate at the stellar surface can be as high as the core mass-loss rate. The extreme mass loss by the core neutrino emission can occur even if the star does not have a hydrogen-rich envelope as long as it is near the Eddington luminosity. Thus, the existence of hydrogen-poor, as well as hydrogen-rich, dense CSM recently observed in some SNe is naturally expected from this mechanism. Even if a star is not close enough to the Eddington luminosity to enhance the mass loss by the neutrino emission shortly before the core collapse, it can still expand because of the reduced gravitational force. This indicates that some stars may have larger radii at the time of the explosion than those expected from the stellar evolution theory.	14	3	1403.2731	We present a novel mechanism for enhancing the mass-loss rates of massive stars shortly before their explosion. The neutrino luminosities of the stellar core of massive stars increase as they get closer to the time of the core collapse. As emitted neutrinos escape freely from the core, the core mass is significantly reduced when the neutrino luminosity is high. If a star is near the Eddington luminosity when the neutrino luminosity is high, the star can exceed the Eddington luminosity because of the core neutrino mass loss. We suggest that the stellar surface mass-loss rates due to the core neutrino emission can be higher than 10<SUP>-4</SUP> M<SUB>⊙</SUB> yr<SUP>-1</SUP> from ~1 year before the core collapse. The mass-loss rates can exceed 10<SUP>-2</SUP> M<SUB>⊙</SUB> yr<SUP>-1</SUP> ~ 10 days before the core collapse. This mass-loss mechanism may be able to explain the enhanced mass loss observed in some supernova progenitors shortly before their explosion. Even if the star is not close enough to the Eddington luminosity to enhance the mass loss, the star can still expand because of the reduced gravitational force. This mechanism can be activated in Wolf-Rayet stars, and it can create the hydrogen-poor, as well as hydrogen-rich, dense circumstellar media observed in some supernovae.	false	[ "massive stars increase", "massive stars", "the core neutrino mass loss", "emitted neutrinos", "the core collapse", "the enhanced mass loss", "Eddington", "the mass loss", "yr", "the stellar core", "The neutrino luminosities", "the neutrino luminosity", "the reduced gravitational force", "Wolf-Rayet stars", "the stellar surface mass-loss rates", "the core neutrino emission", "~1 year", "the Eddington luminosity", "hydrogen-rich, dense circumstellar media", "the core" ]	3.719799	3.93672	-1
745721	[ "Rygl, K. L. J.", "Goedhart, S.", "Polychroni, D.", "Wyrowski, F.", "Motte, F.", "Elia, D.", "Nguyen-Luong, Q.", "Didelon, P.", "Pestalozzi, M.", "Benedettini, M.", "Molinari, S.", "André, Ph.", "Fallscheer, C.", "Gibb, A.", "di Giorgio, A. M.", "Hill, T.", "Könyves, V.", "Marston, A.", "Pezzuto, S.", "Rivera-Ingraham, A.", "Schisano, E.", "Schneider, N.", "Spinoglio, L.", "Ward-Thompson, D.", "White, G. J." ]	2014MNRAS.440..427R	[ "A Herschel and BIMA study of the sequential star formation near the W 48A H II region" ]	7	[ "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy; European Space Research and Technology Centre (ESA-ESTEC), Keplerlaan 1, PO Box 299, NL-2200 AG Noordwijk, the Netherlands", "SKA South Africa, 3rd Floor, The Park, Park Rd, Pinelands 7405, South Africa; Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany; Hartebeesthoek Radio Astronomy Observatory, PO Box 443, Krugersdorp 1740, South Africa", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy; Department of Astrophysics, Astronomy and Mechanics, Faculty of Physics, University of Athens, Panepistimiopolis, GR-15784 Zografos, Athens, Greece", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany", "Laboratoire AIM Paris-Saclay, CEA/IRFU CNRS/INSU Université Paris Diderot, F-91191 Gif-sur-Yvette, France", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy", "Canadian Institute for Theoretical Astrophysics (CITA), University of Toronto, 60 St. George Street, Toronto, ON M5S 3H8, Canada", "Laboratoire AIM Paris-Saclay, CEA/IRFU CNRS/INSU Université Paris Diderot, F-91191 Gif-sur-Yvette, France", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy", "Laboratoire AIM Paris-Saclay, CEA/IRFU CNRS/INSU Université Paris Diderot, F-91191 Gif-sur-Yvette, France", "Department of Physics & Astronomy, University of Victoria, PO Box 355, STN CSC, Victoria, BC V8W 3P6, Canada; National Research Council Canada, 5071 West Saanich Road, Victoria, BC V9E 2E7, Canada", "Department of Physics and Astronomy, University of British Columbia, 6224 Agricultural Road, Vancouver, BC V6T 1Z1, Canada", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy", "Laboratoire AIM Paris-Saclay, CEA/IRFU CNRS/INSU Université Paris Diderot, F-91191 Gif-sur-Yvette, France; Joint ALMA Observatory, Alonso de Cordova 3107, Vitacura 763-0355, Santiago, Chile", "Laboratoire AIM Paris-Saclay, CEA/IRFU CNRS/INSU Université Paris Diderot, F-91191 Gif-sur-Yvette, France; Institut d'Astrophysique Spatiale, UMR8617, CNRS/Université Paris-Sud 11, F-91405 Orsay, France", "ESA/ESAC, PO Box 78, E-28691 Villanueva de la Canada, Madrid, Spain", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy", "Université de Toulouse, UPS-OMP, IRAP, F-31028 Toulouse, France; CNRS, IRAP, 9 Av. colonel Roche, BP 44346, F-31028 Toulouse cedex 4, France", "Infrared Processing and Analysis Center, Institute of Technology, Pasadena, CA 91125, USA", "OASU/LAB-UMR5804, CNRS, Université Bordeaux 1, F-33270 Floirac, France", "Istituto di Astrofisica e Planetologia Spaziali (INAF-IAPS), Via del Fosso del Cavaliere 100, I-00133 Roma, Italy", "Jeremiah Horrocks Institute, University of Central Lancashire, Preston, Lancashire PR1 2HE, UK", "Space Science and Technology Department, STFC Rutherford Appleton Laboratory, Chilton, Didcot, Oxfordshire OX11 0QX, UK; Department of Physics and Astronomy, The Open University, Walton Hall, Milton Keynes MK7 6AA, UK" ]	[ "2016A&A...591A..19P", "2017A&A...607A..22R", "2017MNRAS.472.2990L", "2019AJ....158...18H", "2019ApJ...884....4K", "2020PASJ...72...54C", "2024MNRAS.527.5049L" ]	[ "astronomy" ]	16	[ "stars: formation", "ISM: clouds", "dust", "extinction", "H II regions", "ISM: individual objects: W 48A", "ISM: molecules", "Astrophysics - Astrophysics of Galaxies" ]	[ "1964ARA&A...2..213B", "1973AJ.....78..929P", "1977ApJ...211L..39B", "1977ApJ...214..725E", "1981ApJ...250..341K", "1983QJRAS..24..267H", "1988A&A...190..289P", "1988ApJ...330..809R", "1989ApJ...343..222T", "1989ApJS...69..831W", "1991A&A...246..249P", "1991ApJ...380L..75M", "1991ApJS...76..617T", "1992A&A...253..541C", "1992ApJ...399L..71C", "1993A&A...276..489O", "1994ApJ...426..249O", "1994ApJS...91..659K", "1994MNRAS.268..291W", "1995A&A...294..825V", "1995ASPC...77..433S", "1996A&A...308..573M", "1996A&A...309..827S", "1996A&A...315..565O", "1996A&AS..120..283H", "1996ApJ...469..238M", "1996ApJ...470..528A", "1996ApJ...472..225S", "1996ApJ...473L.131X", "1997A&A...323..931W", "1998A&A...336..339M", "1998AJ....115.1693C", "1998JQSRT..60..883P", "1998MNRAS.301..640W", "1998SPIE.3357..319K", "1999ApJ...517..209G", "1999MNRAS.303..659H", "2000A&A...355..617M", "2000A&A...361..327V", "2000A&A...362..723C", "2000A&A...362.1093M", "2001MNRAS.326..805C", "2002ApJ...566..945B", "2002ApJ...566..982Z", "2002IAUS..206..167V", "2003A&A...403.1095M", "2003ASSL..285..109G", "2003ApJ...593L..51C", "2003ApJ...596..344C", "2004A&A...414.1017T", "2004A&A...421..195C", "2004ApJ...606..917R", "2004MNRAS.355..553G", "2005A&A...433..565D", "2005A&A...436.1049M", "2005ApJ...626..253R", "2005ApJS..157..279A", "2005sf2a.conf..721P", "2006A&A...447..221B", "2006A&A...450..569P", "2006ApJ...645..337H", "2006ApJS..165..283A", "2006MNRAS.367..553P", "2006PASJ...58L..29M", "2007A&A...465..219P", "2007A&A...476.1243M", "2007ARA&A..45..339B", "2007PASJ...59..199O", "2008A&A...481..345M", "2008A&A...487..993K", "2008AN....329..295C", "2008ASPC..387..452H", "2008ApJ...685.1005Q", "2008ApJ...686..384D", "2008ApJS..175..277D", "2009A&A...499..811L", "2009ApJ...693..419Z", "2009ApJ...696..268Z", "2009PASP..121...76C", "2010A&A...510A...5P", "2010A&A...515A..42R", "2010A&A...517A..56F", "2010A&A...518L...1P", "2010A&A...518L...2P", "2010A&A...518L...3G", "2010A&A...518L..77M", "2010A&A...518L..81Z", "2010A&A...518L..83S", "2010A&A...518L..84H", "2010A&A...518L..88B", "2010A&A...518L..97E", "2010A&A...518L.101Z", "2010A&A...524A..18B", "2010AJ....140.1868W", "2010ASPC..434..139O", "2010MNRAS.401.2219B", "2011A&A...527A.111J", "2011A&A...530A.118P", "2011A&A...530A.133M", "2011A&A...533A..94H", "2011A&A...535A..76N", "2011A&A...535L...8G", "2011A&A...536A..38M", "2011ASPC..442..281J", "2011ApJ...735...64W", "2011ApJ...740L...5C", "2011ApJ...742..105A", "2011MNRAS.414..321D", "2011MNRAS.416.2932T", "2012A&A...538A.137M", "2012A&A...539A.156G", "2012A&A...542A.114H", "2012A&A...542L..18S", "2012A&A...543L...3H", "2012A&A...546A..74D", "2012ApJ...752...55K", "2012ITIP...21.3687P", "2012MNRAS.422.1352D", "2013A&A...549A...5R", "2013A&A...550A..50M", "2013A&A...557A..29C", "2013A&A...558A.125D", "2013A&A...559A.133Y", "2013ApJ...762..120P", "2013ApJ...766...85R", "2013ApJ...767..126S", "2013ApJ...772...45E", "2013ApJ...773..102F", "2013MNRAS.435..524B", "2013PASP..125.1126R", "2014A&A...564A.106T" ]	[ "10.1093/mnras/stu300", "10.48550/arXiv.1403.2426" ]	1403	1403.2426.txt	Sequential star formation in the structure of OB associations was noted already in 1960's (e.g., \citealt{blaauw:1964,elmegreen:1977} and references therein). More recently, also on smaller scales stellar age gradients are observed, such as for bright-rimmed clouds (see e.g.,\citealt{white:1997,thompson:2004aa,matsuyanagi:2006,ogura:2007}) and for star formation on the borders of H{\sc ii} regions (e.g., \citealt{deharveng:2005,zavagno:2010a}). Evidence for sequential star formation should also be present during the earliest stages of star formation. Molecular clouds containing star-forming objects at various stages of evolution are commonly observed: the {\em Spitzer} study of \citet{qiu:2008} shows the distributions of the various evolutionary classes of young stellar objects (YSOs) within one cloud; and molecular line studies of infrared dark and high-extinction clouds find collapsing dense infrared-dark cores next to active infrared-bright star-forming ones (see, e.g. \citealt{palau:2010}, \citealt{rygl:2010b,rygl:2013b}). More recently, a few {\em Herschel} HOBYS studies have investigated star formation near OB clusters and bubble-shaped objects in view of triggered star formation (e.g, RCW\,120: \citealt{zavagno:2010}, Rosette: \citealt{schneider:2010,schneider:2012}, \citealt{tremblin:2014}, W5-E: \citealt{deharveng:2012}, W\,3: \citealt{rivera:2013}, NGC 7538: \citealt{fallscheer:2013}). It has not been proven yet whether all sequential star formation is the consequence of triggering, and evidence to the contrary has been found for the Rosette region through observations (\citealt{schneider:2012, cambresy:2013}) and simulations (e.g., \citealt{dale:2011,dale:2012}). Determining age gradients across a star-forming region can give insight into the star formation history by which one can assess the possibility of triggering and its size scale. In this paper we analysed the age gradient in W\,48A, a region with a trustworthy distance and various stages of star formation, tapping into new dust continuum and spectral line data as well as archival data. W\,48A is part of a group of H{\sc ii} regions (\citealt{onello:1994}) located at a distance of $3.27\pm0.49$\,kpc (maser parallax, \citealt{zhangb:2009}) about 1$\rlap{.}$\degr7 below the Galactic plane. \citet{arthur:2006} describe it as a ``classical" champagne flow plus stellar wind \hiir\ where the photo dissociation region (PDR) is moving with 2.5\,\kms\ into the molecular cloud (\citealt{roshi:2005}). The star-forming region W\,48A is known to host several different stages of high-mass star formation: an ultra compact (UC) H{\sc ii} region (G\,35.20--1.74, \citealt{wood:1989a}) embedded in a larger, almost 2\arcmin$\times$2\arcmin, cometary part (\citealt{wood:1989a,kurtz:1994,onello:1994, roshi:2005}) extending to the southeast, a methanol maser-emitting YSO (\citealt{minier:2000}) and some recently discovered prestellar cores containing cold and deuterated molecules (\citealt{pillai:2011}). W\,48A was observed with the {\em Herschel} Space Observatory (\citealt{pilbratt:2010}) at 70 to 500\,$\mu$m under the {\em Herschel} imaging survey of OB Young Stellar objects programme (HOBYS, \citealt{motte:2010}) as a part of a larger map of the W\,48 molecular cloud complex (see the entire {\em Herschel} maps in \citealt{nguyen:2011}). These observations are a powerful tool to visualise the large scale structure of the cold (160--500\,$\mu$m) and warm (70\,$\mu$m) dust emission. {\em Herschel}'s wavelength coverage is ideal to sample the peak of the spectral energy distribution (SED) of cold and warm (10--30\,K) dusty prestellar material and envelopes around YSOs to obtain their temperatures and masses. Since high-mass stars form deeply embedded in their natal cloud, the bulk of their emission is absorbed by the surrounding envelope and is re-emitted in the infrared to mm regime. The infrared luminosity is therefore important for obtaining the bolometric luminosity and determining the spectral type of the embedded star. With the help of a luminosity-mass ($L/M$) diagram (\citealt{saraceno:1996,molinari:2008,elia:2010,hennemann:2010}) one can characterise the evolutionary stage of the YSO through its bolometric luminosity and envelope mass. The {\em Herschel} maps were combined with BIMA+IRAM 30m interferometric molecular line imaging to provide a sub-arcsecond resolution view into the star-forming region, down to size scales of $\sim$1000\,AU, and with several archival data sets from the VLA and JCMT. Molecular line and radio continuum data offer an alternative YSO age estimation to that from the $L/M$ diagram. In particular, we used the \amm\,(1,\,1) and (2,\,2) lines to trace the earlier stages of star formation, the CO\,(3--2) line to find evidence of outflows, and the \mecn\,(6--5), $K$=1, 2, 3, 4, transitions and 107\,GHz methanol masers to trace more the hot and dense gas around YSOs. Furthermore we used the molecular line data to obtain kinematical information: methanol masers for the small scale kinematics around the star-forming objects, and \c18o\,(1--0) data for the large scale kinematics.	\subsection{W\,48A revisited} \citet{roshi:2005} sketched the W\,48A \hiir\ as a PDR dominated H{\sc ii} region, where the ionising star is moving with respect to the molecular cloud. Their C76$\alpha$ radio recombination line (RRL) measurements showed that the PDR ($V_\mathrm{LSR}$=41.9\,kms) was moving at 2.5$\pm$1\,\kms\ into the molecular cloud for which they took $V_\mathrm{LSR}$=44\,\kms\ (\citealt{churchwell:1992}). The LSR velocities of our \c18o\,(1--0) data show that the molecular cloud covers a range of velocities from $\sim$41 up to 45\,\kms . The LSR velocity of the C76$\alpha$ line falls in the velocity range of the molecular material: the PDR is moving {\em at the same velocity} as that part of the cloud. To understand if the PDR is pressure confined by the molecular material, we use the formulae given in \citet{roshi:2005} and calculate pressure of the molecular cloud. The molecular cloud pressure is given by the sum of the pressure from thermal processes and from turbulence (\citealt{xie:1996}). \citet{roshi:2005} demonstrated that this $P_\mathrm{PDR}=5.3-43.0\times10^{-7}\,\mathrm{dyne\,cm^{-2}}> P_\mathrm{UCHii}=1.3\times10^{-7}$\,dyne\,cm$^{-2}$ and that therefore the \hiir\ is pressure confined by the PDR. Now, the pressure in the molecular material in the ridge is given by \begin{equation} P_\mathrm{mol}= n_\mathrm{H_2} (k_BT_\mathrm{mol} + \mu m_\mathrm{H} dv_\mathrm{mol}^2), \end {equation} where $T_\mathrm{mol}=26\,$K (from the dust temperature map), $\mu=2.8$ and $dv_\mathrm{mol}=$1.0--1.8\,\kms\ based on \c18o\,(1--0) and \amm\,(1,1) line widths resulting in $P_\mathrm{mol}=n_\mathrm{H_2} 1\times10^{-13}$\,dyne\,cm$^{-2}$. Using the \amm\ rotational temperature (22\,K) would yield a similar outcome. Estimates of the volume density based on the column density map depend on the depth of the ridge. Assuming the same depth as width ($\sim$1\,pc), the volume density of the molecular material in the ridge would be $6\times10^4$\,cm$^{-3}$. This would bring the pressure of the molecular material to $6\times10^{-9}$\,dyne\,cm$^{-2}$, much lower than the PDR pressure. To stop the PDR from expanding into the molecular cloud, $P_\mathrm{mol}$ would have to be larger than $P_\mathrm{PDR}$ which would require molecular hydrogen densities of $10^8$\,cm$^{-3}$. However, our hydrogen density estimate for the ridge is several orders of magnitude smaller. Therefore, the PDR is expanding into the cloud and pushing the molecular material outward, maintaining the arc-like structure. To fully appreciate the role of the PDR, high-sensitivity RRL imaging of the W\,48A \hiir\ and its surroundings would be necessary to trace the PDR kinematics. Finally, \citet{roshi:2005} found that the \hiir\ is receding with a few ($\sim$3) \,\kms , which might explain why the arc is receding rather than advancing with respect to the LSR velocity of the ridge. In Fig.~ \ref{fig:sketch}, we modify the view of \citet{roshi:2005} to reflect the results of our observations. %FIG 14 \begin{figure} \includegraphics[angle=0,width=8cm]{hii_dec13_fig14.eps} \caption{\label{fig:sketch} Schematic drawing of the W\,48A star-forming region. The cross marks the location of the \hiir\ (white) with LSR velocity of 47.9\,\kms\ (\citealt{wood:1989a}). The PDR is pink and has an LSR velocity of 41.9\,\kms\ (\citealt{roshi:2005}). The molecular cloud is green and contains a wider range of velocities. The purple arrows in the PDR show that it is expanding. Core H2a is marked in dark blue at the border of the PDR. Light blue marks clump H-3 with the H-3b core inside it in dark blue. The molecular outflow of clump H-3 is shown by a red and blue cone indicating also the velocity range. The \hiir\ is moving away with about 3\,\kms . The champagne flow of the \hiir\ is indicated by dashed yellow arrows. } \end{figure} \subsection{Comparison of age estimates} \label{disc:ages} In Section \ref{sec:compact_sources} we derived the ages of the clumps/cores based on the $L/M$ tracks of \citet{molinari:2008} to be: $\sim$$8\times10^5$\,yr for clump H-1 (\hiir) and core H-2b (maser core), and $\sim$$1.5\times10^5$\,yr for core H-3b (containing the molecular outflow). These ages contain many uncertainties: the uncertainties of the envelope masses and bolometric luminosities, and those from stellar evolution models used for the stellar tracks. It is therefore useful to compare these age estimations with the results from radio continuum data, molecular line data, and the literature. Radio continuum and RRL measurements (\citealt{wood:1989a,onello:1994,roshi:2005}) unmistakably showed that the object matching with clump H-1 is a \hiir . Literature ages of \hiir s are similar to our age estimate: $6\times10^5$\,yr (\citealt{motte:2007}), $\sim10^5$\,yr (\citealt{wood:1989a,akeson:1996}). Apart from water and hydroxyl masers, core H-2a contains several class II methanol maser transitions, indicating the presence of a high-mass star in formation (\citealt{breen:2013}). Methanol masers are thought to be excited before the onset of the H{\sc ii} region and be quenched by the ionising radiation of the star (during the H{\sc ii} region phase), and should hence have younger or similar ages as H{\sc ii} regions, i.e., $\lesssim10^5$\, yr (\citealt{walsh:1998,codella:2000,breen:2010}). The detection of \mecn, a molecule typical of intermediate to high-mass protostellar objects such as hot cores (e.g, \citealt{olmi:1996,araya:2005,purcell:2006}), and the absence of a compact radio source (due to the free-free emission of the ionising star, which is typical of H{\sc ii} regions) also agrees with a younger evolutionary stage than a \hiir\ or a deeply embedded \hiir . Higher sensitivity radio observations could test if the object might possibly be a hyper compact H{\sc ii} region. Core H-3b was found to contain a high-mass young stellar object identified by the massive collimated outflow observed in CO\,(3--2) emission (Section~\ref{sec:outflow}) and HCO$^+$(1--0) line (\citealt{pillai:2011}). We detected \amm\ and \mecn\ emission and obtained the [\mecn]/[\amm] abundance ratio, which can be used as a chemical clock (\citealt{charnley:1992}). The abundance ratio found toward core H-3b was $\sim$1$\times10^{-3}$ (see Tables \ref{tab:ch3cn} and \ref{tab:nh3}), which is similar to the typical hot core abundance ratio of [\mecn]/[\amm] $\sim10^{-3}$ that corresponds to ages of $\sim$6$\times10^4$ yr (\citealt{charnley:1992}). Literature ages for hot cores (high-mass protostellar objects) are between $2\times10^4$\,yr (\citealt{motte:2007}) and $5\times10^4$\,yr (\citealt{cazaux:2003}). These two age estimates are both lower than what we obtained from the $L/M$ diagram. Using several molecular lines and radio continuum data, we could infer that: 1) clump H-1 is indeed the most evolved and oldest object, 2) core H-2a is in an earlier evolutionary stage than clump H-1 and most likely also younger than clump H-1, while the $L/M$ diagram based age estimation found similar ages for these two objects, and 3) that core H-3b is indeed the youngest and less evolved object. Given that all three objects are forming high-mass stars the evolutionary difference is likely to reflect the age difference, since high-mass stars evolve along similar tracks, i.e. change their envelope mass and luminosity on similar timescales. It is thus important to have higher-resolution submm or radio continuum data and molecular line data to complement the age estimations based on infrared continuum data (such as the $L/M$ diagram). %FIG 15 \begin{figure} \includegraphics[angle=-90,width=13cm]{fig15_v3.eps} \caption{\label{fig:overview} {\em Panel ($a$):} Sequential star formation around W\,48A indicated on the PACS 70\,$\mu$m map. {\em Panels ($b$) and ($c$):} {\em Herschel} H$_2$ column density and temperature maps of the W\,48 H{\sc ii} region complex (W\,48A--E). The column density map is overlaid with the 21\,cm contours of the NVSS survey (\citealt{condon:1998}). The Aquila supershell expansion is oriented in the direction of Galactic latitude, $b$, indicated in the bottom panel.} \end{figure} \subsection{Sequential star formation sequence around W\,48A?} The top panel in Fig.~\ref{fig:overview} shows a synthesis of our sequential star formation results in W\,48A. In this study we isolated at least four different stages of star formation and estimated their ages, the \hiir\ (clump H-1), the maser core (core H-2a), the active star-forming YSO with outflow (core H-3b), and the cold streamers in which \citet{pillai:2011} found prestellar cores in NH$_2$D of 50--200\,$M_\odot$. The ages of prestellar objects depend strongly on whether they will form low- or high-mass stars. Assuming that these massive prestellar cores will form high-mass stars, then they should have younger ages than hot cores $<5\times10^4$yr (\citealt{cazaux:2003}). As protostellar heating destroys NH$_2$D, detecting these prestellar cores indicates that no protostars have been formed yet, which was confirmed through the non-detection of protostellar objects in the 70\,$\mu$m PACS map. Interestingly, these four stages of star formation are aligned in a east (old) to west (young) configuration, forming a linear age gradient (Fig.~\ref{fig:overview}, panel ($a$)). Large projection effects can, most likely, be disregarded since the objects are forming within the same molecular cloud, suggesting that this remarkable age arrangement is real and reflects the formation history of the W\,48 molecular cloud. Panels ($b$) and ($c$) in Fig.~\ref{fig:overview} show that the \hiir\ formed on the border of the molecular cloud. Core H-2a could have either formed before the creation of the arc and was swept up by it, or formed in the material collected by the PDR in a collect and collapse mechanism (\citealt{elmegreen:1977,whitworth:1994}). It is unlikely that the \hiir\ triggered the formation of the W\,48A molecular cloud nor the star formation inside it. In the previous section we mentioned the \hiir\ is confined by the PDR (\citealt{roshi:2005}), but that the PDR is slowly expanding into the molecular medium. The distance between \hiir\ and clump H-3 is about 1\,pc. A shock wave propagating at $\sim$0.3\,\kms\ , which is the sound speed of the medium at 25\,K, would take about 3\,Myr. The age clump H-3 is about an order of magnitude younger than that, ruling out the \hiir\ (and the PDR) as the triggering agent. Hence, the W\,48A molecular cloud and its star-forming seeds were assembled before the onset of the \hiir\ in such a way that star formation started in the east and progressed toward the west. The sequential star formation in W\,48A reminds of the age gradient seen in the DR\,21 molecular cloud which was formed by colliding flows, and has a south (old) - north (young) age gradient (see \citealt{csengeri:2011b,hennemann:2012}). Given the location of W\,48 at almost 100\,pc below the Galactic plane, the presence of an age gradient suggests a large scale external force. First of all, the location of the other W\,48 H{\sc ii} regions cannot explain the age gradient (by e.g., triggering) and it seems more likely that all these H{\sc ii} regions belong to one larger star-forming complex. Based on 21\,cm HI line data, \citet{maciejewski:1996} discovered the Aquila supershell, centred at about four degrees below W\,48A, at $l$=35$^\circ$, $b$=--6$^\circ$. This shell is expanding from below into the Galactic plane, creating, in addition to a super bubble, a cone-like shape directed at $l$=$34\rlap{.}^\circ6$, $b$=$-1\rlap{.}^\circ4$, very near ($\sim$40\arcmin\ or $\sim$38\,pc) to the W\,48 H{\sc ii} regions. The 21\,cm HI spectra of the left and centre part of this cone show a clear peak at 40--45\,\kms (Figure 4, \citealt{maciejewski:1996}), which overlaps with the LSR velocity of the molecular material of W\,48A and of the other H{\sc ii} regions in W\,48: 41.2\,\kms\ for W\,48B; 46.7\,\kms\ for W\,48D; 45.5\,\kms\ for W\,48E (\citealt{onello:1994}). It is therefore likely that the initial molecular cloud in which the W\,48 complex of H{\sc ii} regions was formed through the compression of ambient material by the Aquila supershell. \citet{maciejewski:1996} estimated that the total energy of the events creating the Aquila supershell are about $1-5\times10^{52}$\,erg, corresponding to 10--100 supernovae explosions powering the system over $10^7$ yr. The W\,48 H{\sc ii} regions were formed recently given the supershell's age ($\sim2\times10^7$\,yr) and its estimated expansion velocity of about 15\,\kms\ (which is variable depending on the density of the material it passes through). The distance between the centre of the supershell and the W\,48 H{\sc ii} regions is about 4 degrees (230\,pc at a distance of 3.27\,kpc) and at a velocity of 15\,\kms\ it would take about 1.5 Myr for the shell to reach the W\,48 H{\sc ii} regions, leaving ample time for the W\,48 H{\sc ii} regions to evolve into their current state. At the location of the W\,48 H{\sc ii} regions, the supershell's orientation is toward increasing Galactic latitudes. The shell would have therefore first reached the locations of W\,48C and D, then W\,48A and B, and finally W\,48E. Hence, the density structures created by the shell, should be older in the eastern side of Fig~\ref{fig:overview} and younger in the western side. Most of these H{\sc ii} regions are too evolved to have strong far-infrared emission, hence it is difficult to estimate their evolutionary stages from their envelope masses and bolometric luminosities. It is more fruitful to look at the size of the ionised hydrogen region (the Str\"omgren shell) created by the young star, which one can observe through its free-free emission in the centimetre radio continuum. For this purpose we used the 21\,cm maps from the NRAO's NVSS survey (\citealt{condon:1998}). With AIPS task JMFIT we fitted 2D Gaussians to the 21\,cm emission of W\,48A--D obtaining the size, peak flux and integrated flux (Table \ref{tab:21cm}). The ratio of the integrated to peak flux (I/P), listed in the fifth column of Table~\ref{tab:21cm}, is a measure of the source compactness (1.0 for a point source and increasing when the source is more extended). W\,48D seems indeed to be an old H{\sc ii} region, since its radio emission very extended and it has possibly triggered another younger and compact H{\sc ii} region, W\,48C, which is very compact, hence very young, on one of its borders. The radio emission of W\,48B is less extended than W\,48D, but more extended that that of W\,48A, implying an intermediate evolutionary stage. This is consistent with W\,48B already containing a cluster of stars (clearly visible in the near-IR) and being surrounded by a small circular dust arc. Based on the 21\,cm radio continuum we conclude that the H{\sc ii} regions have an evolutionary gradient along Galactic latitude, such as one would expect if these regions were formed by the Aquila supershell. In this discussion we have neglected W\,48E, whose nature, either a H{\sc ii} region or a supernova remnant, is uncertain, since it emits extended weak radio emission, and does not contain many YSOs, nor high dust column densities. \begin{table} \begin{flushleft} \caption{21\,cm emission properties of the W\,48 H{\sc ii} regions} \label{tab:21cm} \begin{tabular}{l c c c c} \hline H{\sc ii} region & Size & Peak Flux & Int. Flux & I/P\\ & (arcsec, arcsec) & (Jy~beam$^{-1}$) & (Jy) &\\ \hline W\,48A & 74, 72 & 4.9 & 12.8 & 2.6\\ W\,48B & 118, 84 & 0.12 & 0.6 & 5.0\\ W\,48C & 50, 49 & 0.16 & 0.2 & 1.3\\ W\,48D & 221, 169 & 0.097 & 1.8 & 185.5\\ \hline \end{tabular} NOTES. Columns are (from left to right): name; size in major and minor axis at FWHM; Peak flux; Integrated Flux; ratio of Integrated to Peak flux. \end{flushleft} \end{table} In addition to the W\,48 H{\sc ii} regions, IRAS 18586+0106, also known as Mol\,87, is located between W\,48A and W\,48E. IRAS 18586 contains a few massive star-forming clumps (\citealt{beltran:2006aa}), but no centimetre continuum emission, which would indicate free-free emission of young stars, could be associated with them (\citealt{molinari:1998a}), nor methanol or water masers (\citealt{fontani:2010}). With its \amm\,(1,1) LSR velocity of 38\,\kms\ (\citealt{molinari:1996}), IRAS 18586 is very likely to be a part of the W\,48 complex. In Appendix~\ref{sec:iras} we analysed the infrared emission of the two clumps found in the {\em Herschel} maps. Source A, which coincides with the IRAS-coordinates centre, is the most evolved source with an age of $\sim$5$\times$10$^5$\,yr forming an intermediate to high-mass star. The second source (B) is colder, more massive and younger. For the beginning of star formation in IRAS 18586, we only need to take the estimate of the oldest object into consideration, which is $\sim$5$\times$10$^5$\,yr. This is younger than most evolved object in W48\,A (the \hiir), and would agree with the predicted age gradient if all these star-forming regions were formed by the cone of the Aquila supershell. In addition, the {\em Herschel} column density and temperature maps (Fig.~\ref{fig:overview}) revealed that there exists a large reservoir of dense and cold gas west of W\,48A, implying that this region has a high potential to form stars. The morphology of this dense gas suggests that it has not been strongly affected by H{\sc ii} regions or stellar winds, that are known to shape the material into bubbles and shells. East of W\,48A, there is much less dense material, and when present it is shaped in shells, bubbles and ridges by H{\sc ii} regions. Hence, the large scale {\em Herschel} data support the hypothesis of an east-west evolutionary gradient across the W\,48 H{\sc ii} regions, in a roughly similar orientation as the age gradient found around W\,48A.	14	3	1403.2426	We present the results of Herschel HOBYS (Herschel imaging survey of OB Young Stellar objects) photometric mapping combined with Berkeley Illinois Maryland Association (BIMA) observations and additional archival data, and perform an in-depth study of the evolutionary phases of the star-forming clumps in W 48A and their surroundings. Age estimates for the compact sources were derived from bolometric luminosities and envelope masses, which were obtained from the dust continuum emission, and agree within an order of magnitude with age estimates from molecular line and radio data. The clumps in W 48A are linearly aligned by age (east-old to west-young): we find a ultra-compact (UC) H II region, a young stellar object (YSO) with class II methanol maser emission, a YSO with a massive outflow and finally the NH<SUB>2</SUB>D prestellar cores from Pillai et al. This remarkable positioning reflects the (star) formation history of the region. We find that it is unlikely that the star formation in the W 48A molecular cloud was triggered by the UC H II region and discuss the Aquila supershell expansion as a major influence on the evolution of W 48A. We conclude that the combination of Herschel continuum data with interferometric molecular line and radio continuum data is important to derive trustworthy age estimates and interpret the origin of large-scale structures through kinematic information.	false	[ "continuum data", "additional archival data", "trustworthy age estimates", "UC H II", "continuum emission", "interferometric molecular line", "class II methanol maser emission", "age estimates", "data", "W 48A.", "et al", "kinematic information", "Berkeley Illinois Maryland Association", "OB Young Stellar objects", "molecular line and radio data", "age", "II", "Herschel imaging survey", "the UC H II region", "Herschel HOBYS" ]	10.924259	10.390296	183
745536	[ "Mattsson, Lars", "De Cia, Annalisa", "Andersen, Anja C.", "Zafar, Tayyaba" ]	2014MNRAS.440.1562M	[ "On the (in)variance of the dust-to-metals ratio in galaxies" ]	48	[ "Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 Copenhagen Ø, Denmark; Nordita, KTH Royal Institute of Technology and Stockholm University, Roslagstullsbacken 23, SE-106 91 Stockholm, Sweden", "Department of Particle Physics and Astrophysics, Faculty of Physics, Weizmann Institute of Science, Rehovot 76100, Israel", "Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 Copenhagen Ø, Denmark", "European Southern Observatory, Karl-Schwarzschildstrasse 2, D-85748 Garching bei München, Germany" ]	[ "2014MNRAS.443..522Y", "2014MNRAS.443.1704H", "2014MNRAS.444..797M", "2015A&A...582A.121R", "2015ApJ...809..147H", "2015ApJ...810...70Z", "2015MNRAS.447.2937H", "2015MNRAS.449.3274F", "2015MNRAS.451..167F", "2015MNRAS.454.2381G", "2015arXiv150504758M", "2016A&A...596A..97D", "2016ApJ...822....9H", "2016MNRAS.457.3775M", "2016MNRAS.459.1646C", "2016P&SS..133..107M", "2016PASJ...68...94H", "2017A&A...599A..24W", "2017ApJ...835..154H", "2017ApJ...846..151H", "2017MNRAS.465...54C", "2017MNRAS.470..771T", "2018A&A...613L...2D", "2018MNRAS.475.3883J", "2018MNRAS.478.4905A", "2018MNRAS.480..108Z", "2019A&A...624A..80N", "2019MNRAS.482.2731Z", "2019MNRAS.489.1397L", "2020MNRAS.491.3937T", "2020MNRAS.491.4334M", "2020MNRAS.499.6035M", "2021A&A...656A..65M", "2021EPJP..136..339P", "2021MNRAS.502...15H", "2021MNRAS.502.4723Z", "2022A&A...659A..46B", "2022MNRAS.509.3218K", "2022MNRAS.509.5771H", "2022MNRAS.515..320N", "2022MNRAS.517.2076M", "2023A&A...679A..91H", "2023ApJ...951...66W", "2023MNRAS.518.4214F", "2024A&A...681A..64K", "2024A&A...683A.216D", "2024A&A...685A.103V", "2024arXiv240107963D" ]	[ "astronomy" ]	6	[ "stars: AGB and post-AGB", "supernovae: general", "dust", "extinction", "galaxies: evolution", "galaxies: spiral", "Astrophysics - Astrophysics of Galaxies" ]	[ "1982A&A...115..373V", "1989IAUS..135..431M", "1990A&A...236..237I", "1996ApJ...469..740J", "1998AJ....116..748G", "1998ApJ...493..583V", "1998ApJ...501..643D", "1998MNRAS.301..569L", "2003PASJ...55..901I", "2004A&A...421..479V", "2004ApJ...614..796S", "2006ApJ...648..435N", "2007ApJ...662..927D", "2007ApJ...663..866D", "2007MNRAS.378..973B", "2009ApJ...692..677D", "2009MNRAS.394.1061H", "2010A&A...515A..68M", "2011EP&S...63.1027I", "2011MNRAS.414..781M", "2011MNRAS.416.1340H", "2011MNRAS.416.1916V", "2011Sci...333.1258M", "2012ApJ...752..112H", "2012ApJ...760...96G", "2012MNRAS.423...26M", "2012MNRAS.423...38M", "2012MNRAS.424.2345V", "2012MNRAS.424L..34K", "2013A&A...551A..25P", "2013A&A...560A..26Z", "2013A&A...560A..88D", "2013ARA&A..51..457N", "2013ApJ...766...59G", "2013ApJ...774....8T", "2013EP&S...65..213A", "2013MNRAS.432..637A", "2013MNRAS.432.1217P", "2013MNRAS.434..451L", "2013MNRAS.436.1238K", "2013arXiv1306.0008C", "2014A&A...563A..31R", "2014Natur.505..186F" ]	[ "10.1093/mnras/stu370", "10.48550/arXiv.1403.0502" ]	1403	1403.0502_arXiv.txt	The variation of the overall dust-to-metals ratios between galaxies of vastly different morphology, ages and metallicities appears surprisingly small in many cases, with a mean value close to the Galactic ratio ($\sim 0.5$). The relatively tight correlation between the dust-to-gas ratio and the metallicity (yielding an almost invariant dust-to-metals ratio) in the Local Group galaxies has been known for quite a while \citep[see][]{Viallefond82,Issa90,Whittet91}. Indirect evidence for a `universal' mean value is also provided by the almost linear relation between $B$-band optical depth and stellar surface density in spiral galaxies \citep{Grootes13}. But recent results based on gamma-ray burst (GRB) afterglows, quasar foreground damped Ly$\alpha$-absorbers \citep[DLAs;][]{Zafar13} and distant lens galaxies \citep[see, e.g.,][]{Dai09,Chen13} now seem to extend this correlation beyond the local Universe and down to metallicties just a few percent of the solar value. \citet{Zafar13} argue this can only be explained by either rapid dust enrichment by supernovae or very rapid interstellar grain growth by accretion of metals. However, there can be significant variations {\it within} a galaxy \citep[see, e.g.,][]{Mattsson12,Mattsson12b}, although the existence of dust-to-metals gradients is somewhat difficult to establish observationally with reliable independent methods. If, on the other hand, the dust-to-metals ratio does not vary much at all, in any environment, one may assume dust grains as well as atomic metals are mainly produced by stars. Recent findings of large amounts of cold dust in supernova (SN) remnants \citep{Matsuura11,Gomez12} seem to support this hypothesis, although the exact numbers can be disputed \citep{Temim13,Mattsson13a,Mattsson13b}. In other words: the overall picture is not consistent. A new study by \citet{DeCia13} seems to confirm the rising trend with metallicity of the dust-to-metals ratio in quasar DLAs found in previous studies \citep{Vladilo98, Vladilo04}. Furthermore, \citet[][see also Herrera-Camus et al. 2012]{Fisher13} derived a dust mass in the local starburst dwarf I Zw 18, as well as a high-redshift object of similar character, which clearly indicate a dust-to-metals ratio below the Galactic value. These results, together with the likely existence of dust-to-metals gradients along galaxy discs \citep{Mattsson12b}, suggest the variance (or invariance) of the dust-to-metals ratio may depend on the environment. In such case, there may exist an equilibrium mechanism that keeps the dust-to-metals ratio close to constant if certain conditions are fulfilled, while a metallicity dependence may occur as a result of deviations from those conditions in other environments. Recently, \citet{Kuo13} have tried to alleviate the tension between the results from the GRB afterglows of \citet{Zafar13} and other data (for local dwarf galaxies) by fine-tuning the parameters of their standard galactic dust evolution model including grain growth \citep{Hirashita11, Kuo12}. What they suggest is a quite reasonable compromise, but a truly convincing explanation of the different trends (constant and rising dust-to-metals ratio) would require some modification of the dust-formation scenario. In particular, a model in which inherent properties of a galaxy more or less uniquely determines its dust-to-metals ratio would be desirable. Even if the models by \citet{Kuo13} are marginally consistent with the data they compare with, there is obviously still some tension between models and observations. The new results by \citet{DeCia13} only act as to emphasise this. In the standard picture of production and destruction of cosmic dust one is faced with the following two problems: (1) in metal-poor environments dust is only supplied by stars as the interstellar density of metals is too low for efficient grain growth, but still being destroyed by SN shockwaves (albeit with a relatively low efficiency); (2) to compensate the destruction of dust grains, which eventually becomes efficient, with grain growth requires that one pushes the boundaries of the model, i.e., to obtain a sufficiently short grain-growth timescale, one is forced to accept a very large span of gas densities (several orders of magnitude) and a very low star-formation efficiency. These problems are discussed in more detail by \citet{Kuo13}. In this paper we investigate a scenario for the evolution of the galactic dust component where destruction of grains due to sputtering in SN shockwaves is roughly balanced by grain growth by accretion of molecular gas. This idea has also been put forth in other studies to improve models of the build-up of dust in the local as well as distant (early) Universe \citep[see, e.g.,][]{Inoue11,Mattsson11,Valiante11}, but here we take it one step further and consider a model where there can be an {\it exact} balance. Given a constant ratio of the effective dust yield and the total metal yield for a generation of stars, such a scenario will lead to an invariant dust-to-metals ratio. We continue by discussing the possibility that young undeveloped (low metallicity) systems may have a different yield ratio due to different dust yields for individual stars (e.g., the expected metallicity dependence).	Several observational studies suggest a surprisingly small variation of the dust-to-metals ratio in vastly different environments. It is worth stressing that the `trivial solution' to the problem, i.e., adopting a (constant) yield ratio of $y_{\rm d}/y_Z \sim 0.5$, works for any model where there is a replenishment mechanism to counteract possible dust destruction \citep[such as the model used by][for example]{Kuo13}. But other observational evidence also suggest there is a significant variation of the dust-to-metals ratio between different environments, and an invariant dust-to-metals ratio is problematic also in the sense that it requires fine-tuning and is pushing the limits of the `standard models' of dust evolution in galaxies to explain all data \citep{Kuo13}. We find that a reasonable way to resolve this apparent contradiction, and avoiding fine-tuning and extreme model parameters, is to assume that stellar dust production can be efficient, but that interstellar dust growth is equally important and act as a replenishment mechanism which can almost exactly counteract the dust destruction in the ISM. In this scenario, the ratio of the effective (stellar) dust and metal yields is not likely a universal constant and may change due to some metallicity-dependence of the stellar dust yield. We propose the existence of a critical stellar metallicity above which nucleation and condensation of dust in stars can be efficient. We conclude that destruction and growth of grains in the ISM likely strives towards an equilibrium state, which mimics the general behaviour of the case of pure stellar dust production (and no destruction of grains). This explains the relatively small variation of the dust-to-metals ratio seen in several observational studies of local galaxies, but allows also for a significantly lower ratio at low metallicity if the effective stellar dust yield can vary with metallicity. The suggested scenario has important implications for the rapid build-up of large dust masses at high redshifts. Instead of requiring an extreme efficiency of dust formation in massive stars (SNe) as suggested by, e.g., \citet{Dwek07}, the large dust masses seen in the quasar-host galaxy SDSS J1148+5251 (and other objects at high redshifts), follows naturally from the rapid production of metals that is expected in a massive starburst. Just as \citet{Valiante11} we are led to conclude that, though massive stars must produce significant amounts of dust, dust masses of the order $10^8-10^9\,M_\odot$ (as in SDSS J1148+5251) are not likely a result of stellar dust sources only (as a consequence of interstellar dust destruction) and the resultant dust component must therefore be dominated by grain growth in molecular clouds.	14	3	1403.0502	Recent works have demonstrated a surprisingly small variation of the dust-to-metals ratio in different environments and a correlation between dust extinction and the density of stars. Naively, one would interpret these findings as strong evidence of cosmic dust being produced mainly by stars. But other observational evidence suggest that there is a significant variation of the dust-to-metals ratio with metallicity. As we demonstrate in this paper, a simple star-dust scenario is problematic also in the sense that it requires that destruction of dust in the interstellar medium (e.g. due to passage of supernova shocks) must be highly inefficient. We suggest a model where stellar dust production is indeed efficient, but where interstellar dust growth is equally important and acts as a replenishment mechanism which can counteract the effects of dust destruction. This model appears to resolve the seemingly contradictive observations, given that the ratio of the effective (stellar) dust and metal yields is not universal and thus may change from one environment to another, depending on metallicity.	false	[ "dust destruction", "interstellar dust growth", "dust extinction", "dust", "cosmic dust", "stellar dust production", "stars", "supernova shocks", "metallicity", "different environments", "a simple star-dust scenario", "metals", "passage", "the interstellar medium", "strong evidence", "the effective (stellar) dust and metal yields", "other observational evidence", "a replenishment mechanism", "that destruction", "Recent works" ]	12.4462	8.151728	187
437721	[ "Samal, M. R.", "Zavagno, A.", "Deharveng, L.", "Molinari, S.", "Ojha, D. K.", "Paradis, D.", "Tigé, J.", "Pandey, A. K.", "Russeil, D." ]	2014A&A...566A.122S	[ "The molecular complex associated with the Galactic H II region Sh2-90: a possible site of triggered star formation" ]	50	[ "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "INAF - Instituto Fisica Spazio Interplanetario, via Fosso del Cavaliere 100, 00133, Roma, Italy", "Department of Astronomy and Astrophysics, Tata Institute of Fundamental Research, Homi Bhabha Road, 400 005, Mumbai, India", "Université de Toulouse, UPS-OMP, IRAP, 31 400, Toulouse, France; CNRS; IRAP, 9 Av. du Colonel Roche, BP 44346, 31028, Toulouse Cedex 4, France", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Aryabhatta Research Institute of Observational Sciences, 263 129, Nainital, India", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France" ]	[ "2014ApJ...797...40Y", "2015A&A...576A.110N", "2015A&A...581A...5S", "2015A&A...582A...2D", "2015ARep...59..360P", "2015ApJ...798...30L", "2015ApJ...800..101A", "2015ApJ...809..154H", "2015MNRAS.450.1199D", "2015MNRAS.454.3597D", "2016A&A...585A..30C", "2016A&A...592A..77B", "2016AJ....152..117Y", "2016ApJ...818...95L", "2016ApJ...822...49J", "2016ApJ...830...57G", "2016JApA...37...38M", "2016MNRAS.458.4222Z", "2016arXiv161102661C", "2017A&A...600A..93F", "2017A&A...602A..95L", "2017A&A...605A..35P", "2017A&A...606A...8D", "2017ApJ...836...98J", "2017ApJ...849..140X", "2017MNRAS.468.2684P", "2017RMxAA..53...79C", "2018A&A...617A..67S", "2018A&A...618A..67C", "2018AJ....155...44P", "2018ApJ...861..117D", "2018ApJ...864..154D", "2018MNRAS.477.4577S", "2018NewA...63...27V", "2019AJ....157..112P", "2019ApJ...884...77B", "2019MNRAS.487.2881S", "2019MNRAS.488.4753D", "2020A&A...637A..40Z", "2020MNRAS.496..870A", "2020MNRAS.496.1051A", "2020MNRAS.499..668G", "2021A&A...646A.103D", "2021ApJ...923..198D", "2021ApJS..254...33K", "2021MNRAS.500.3123D", "2021MNRAS.504.2557D", "2022MNRAS.510.4436M", "2023ApJ...948....7D", "2023ApJ...953..145V" ]	[ "astronomy" ]	8	[ "HII regions", "stars: formation", "stars: protostars", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]	[ "1959ApJS....4..257S", "1975A&A....42..273G", "1976A&AS...25..179C", "1977A&A....60..233I", "1977ApJ...214..725E", "1977ApJ...218..148M", "1978ApJ...224..132B", "1979A&A....71...59T", "1982A&AS...47....1H", "1982ApJS...49..183B", "1983A&A...124....1L", "1983QJRAS..24..267H", "1984ApJ...281..624S", "1985ApJ...288..618R", "1986A&A...167..157T", "1986A&A...169..281C", "1987PASP...99..191S", "1988PASP..100.1134B", "1990ApJ...354..529B", "1991CSci...60...95S", "1992AJ....104..340L", "1993A&A...275...67B", "1994ApJS...91..659K", "1994MNRAS.268..291W", "1995A&A...302..521C", "1996A&A...311..858B", "1996A&AS..117..393B", "1997AJ....114..288M", "1997pism.book.....D", "1998A&A...337..403B", "1998A&AS..132..211H", "2000A&A...358..593S", "2000prpl.conf...59A", "2001ApJ...548..296W", "2001ChJAA...1...60X", "2002ApJ...566..945B", "2002ApJS..141..157Y", "2002MNRAS.336..621L", "2002MNRAS.337.1309S", "2003AAS...203.5708E", "2003PASP..115..953B", "2003yCat.2246....0C", "2004ApJ...613..986P", "2004ApJS..154...10F", "2004ApJS..154..363A", "2004RvMP...76..125M", "2004SPIE.5492..978P", "2005A&A...433..565D", "2005A&A...436.1049M", "2005ApJ...629..881H", "2005IAUS..227..206W", "2006A&A...446..171Z", "2006A&A...450..607W", "2006A&A...450..625U", "2006A&A...457..637M", "2006ApJ...649..759C", "2006ApJ...651..502P", "2006ApJ...653.1226Q", "2007A&A...461..197W", "2007A&A...472..835Z", "2007ARA&A..45..339B", "2007ApJ...663.1069F", "2007ApJS..169..328R", "2007MNRAS.379.1599L", "2008A&A...481..345M", "2008A&A...490L..27A", "2008ASPC..387..290R", "2008ApJ...672..410L", "2009A&A...497..649B", "2009PASP..121...76C", "2010A&A...518L...2P", "2010A&A...518L...3G", "2010A&A...518L.100M", "2010A&A...523A...6D", "2010ApJ...709..791B", "2010ApJ...714.1015S", "2010ApJ...723L...7K", "2010ApJ...724L..44C", "2010PASP..122..314M", "2011A&A...526A.151R", "2011A&A...529L...6A", "2011A&A...530A.133M", "2011A&A...533A..94H", "2011ApJ...741..110D", "2011MNRAS.414..321D", "2011MNRAS.416.2932T", "2012A&A...539A.156G", "2012A&A...540L..11S", "2012A&A...541A.132P", "2012A&A...542A..10A", "2012A&A...543L...3H", "2012A&A...546A..74D", "2012ApJ...755...20S", "2012ApJ...755...71K", "2012MNRAS.421..408T", "2012MNRAS.427.2852D", "2013A&A...552A..14O", "2013A&A...554A...6R", "2013ARep...57..573P", "2014yCat.2328....0C" ]	[ "10.1051/0004-6361/201321794", "10.48550/arXiv.1403.7925" ]	1403	1403.7925_arXiv.txt	} There are several ways the energy inputs from the OB stars can modify the physical environment and chemistry of the host complex in which they reside (e.g., McKee \& Ostriker 1977), and therefore can trigger the formation of a new generation of stars in the complex (e.g., Elmegreen \& Lada 1977; Bertoldi \& McKee 1990). Recent observational studies of bubbles associated with \hii regions (e.g., Deharveng et al. 2010), suggest that their expansion possibly triggers 14\% to 30\% of the star formation in our Galaxy (e.g., Deharveng et al. 2010; Thompson et al. 2012; Kendrew et al. 2012). These observational results have revealed the importance of OB stars on star formation activity on a Galactic scale. The detailed studies of individual objects (e.g., Deharveng et al. 2006, 2008; Urquhart et al. 2006; Zavagno et al. 2006, 2007), however, showed that the complex environments around \hii regions makes determining the exact process of star formation difficult and that, in general, this process is complicated. Now the far-infrared (FIR) observations provided by Hi-GAL ({\it Herschel} Infrared Galactic Plane Survey; Molinari et al. 2010a) using the {\it Herschel} Space Observatory have the ability to image a large area of a cloud complex with unprecedented sensitivity, thus allowing us to improve our understanding of how OB stars interact with the local interstellar medium (ISM), and process cold gas to induce new star formation. The recent studies based on {\it Herschel} observations revealed that star-forming regions (SFRs) are composed of very complex and filamentary clouds, with ongoing star formation at various locations (e.g., Hill et al. 2011; Giannini et al. 2012; Hennemann et al. 2012, Deharveng et al. 2012; Schneider et al. 2012; Preibisch et al. 2013; Roccatagliata et al. 2013). Demonstrating the existence of triggered star formation in extended clumpy clouds by internal feedback sources is difficult, because they may be forming new stars in various ways. For example, such clouds can form stars spontaneously governed by the physical condition and the evolution of the original cloud or the collapse of high-density structures generated by the large-scale supersonic turbulence of the the ISM (e.g., see Mac Low \& Klessen 2004). Thus, understanding of the physical connection and interaction of bubbles/\hii regions with the cold ISM, and their association with stellar/proto-stellar content is central to obtaining a better picture of star formation around \hii regions. In this context, we present here a multiwavelength investigation of the Sh2-90 \hii complex (briefly described in Sect. 2) in order to decipher its star formation activity. In the present work, we analyzed the distributions of the ionized and cold neutral ISM, with radio continuum observations at low frequencies (610 and 1280 MHz) and dust continuum emission with {\it Herschel} in the wavelength range 70-500 $\mu$m. We explore the stellar and proto-stellar components of the complex, using high-resolution $JHK$ observations coupled with archival {\it Spitzer}-IRAC observations. We have organized this work as follows. Section 2 presents the Galactic \hii region Sh2-90. In Sect. 3, we describe the observations and the reduction procedures. Section 4 describes the \hii region (adopted distance, general morphology, properties of ionizing gas, and exciting source). Section 5 deals with the distribution and physical condition of the cold ISM, and the properties of compact dusty clumps. In Sect. 6, we describe the identification and classification of young stellar objects (YSOs), their nature and distribution. The kinematics of ionized and molecular gas is presented in Sect. 7. Section 8 is devoted to the general discussion and our understanding of star formation in the Sh2-90 complex. We present the main conclusions in Sect. 9. \begin{table} \centering \scriptsize \caption{ IRAS point sources towards the Sh2-90 complex} \begin{tabular}{cccccccc} \hline\hline Name & RA & Dec & F$_{12}$ & F$_{25}$ & F$_{60}$ & F$_{100}$ & L \\ & deg (J2000) & deg (J2000) & Jy & Jy & Jy & Jy & 10$^3$ \lsuns \\ \hline IRAS 19473+2638 & 297.341461 & 26.775908 & 6.02 & 14.25 & 111.90 & 1514.00 & 5.7 \\ IRAS 19474+2637 & 297.385773 & 26.753876 & 4.20 & 14.75 & 92.02 & 3193.00 & 10.4\\ IRAS 19473+2647 & 297.358250 & 26.920030 & 2.62 & 1.40 & 3.31 & 49.50 & 0.3\\ IRAS 19471+2641 & 297.291168 & 26.814301 & 7.68 & 29.29 & 285.80 & 1673.00 & 7.7\\ IRAS 19470+2643 & 297.287231 & 26.849007 & 2.93 & 44.75 & 330.60 & 1514.00 & 7.6 \\ \hline \end{tabular} \end{table} \begin{figure} \center \includegraphics[width=13cm,height=13cm ]{Fig1.jpeg} \caption{Colour-composite image of the Sh2-90 complex. {\it Spitzer}-IRAC 8.0 \mum (red) and 3.6 \mum (green) images have been combined with the DSS2 R-band (blue) image. The different sources associated with the region (see the text) are marked. The field size is 12$\farcm$0 (E-W) $\times$ 12$\farcm$0 (N-S), centered at $\alpha_{2000} = 19^{\rm h}49^{\rm m}18^{\rm s}$, $\delta_{2000} = +26^{\circ}49^{\prime}29^{\prime\prime}$. North is up and east is left.} \label{fig1} \end{figure}	} In this paper we have presented {\it Herschel} and radio continuum imaging of the Sh2-90 complex as well as {\it Spitzer}-IRAC and deep NIR imaging to explore its stellar and interstellar content, and star formation history. Based on these observations our results can be summarized as follows: 1. The Sh2-90 complex consists of two bubbles (N133 and N132) and a few IRAS sources at various locations. N133 is a large bubble outlined by {\it Spitzer} 8.0 $\mu$m emission and together with Sh2-90, encloses the main \hii region of the complex. It is an evolved \hii region of diameter $\sim$ 3.2 pc with an rms electron density $\sim$ 144 cm$^{-3}$, and an ionized mass $\sim$ 55 \msun. Our NIR photometry of the sources inside the bubble reveals the presence of a loose cluster, with the most massive member of which is a O8-O9 V star that is responsible for the ionization of N133. 2. The column density and temperature maps constructed from the {\it Herschel} observations suggest that Sh2-90 is part of a massive ($\ge$ 10$^{4}$ \msun) elongated cloud of column density $\geq$ 3 $\times$ 10$^{21}$ cm$^{-2}$. We observed that neutral collected material of mass $\ge$ 637 $\msun$ is present in a shell surrounding the \hii region. Nine clumps are detected in the complex, among which seven (mass range 8-125 $\msun$) are located at the periphery or in the direction of Sh2-90. Four of them are co-spatial with B-type stars and a compact \hii region. The velocity information of the clumps derived from CO (J=3-2) data cubes suggests that most of them are likely to be associated with the Sh2-90 complex. 3. Using the IRAC and NIR CC diagrams, we identified 129 likely YSO candidates of masses in the range of 0.2-3 $\msun$, which includes 21 Class I, 34 Class II, and 74 NIR-excess YSOs. The photometric measurements of these YSOs are available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/gcat?J/A+A. We identified four Class 0 MYSO in their main accretion phase. We observed that the spatial distribution of the candidate YSOs follows the distribution of the high column density matter, with YSOs clustering at various locations, indicating that recent star formation is going on at multiple sites. 4. We find the possible existence of two generation of massive to intermediate-mass star formation in the complex; one is in the immediate vicinity of the Sh2-90 \hii region, in the form of NIR/optical point sources responsible for the excitation of compact IR-blobs and a compact \hii region, and the other in the form of young ($\sim$ a few $\times$ 10$^{4}$ yr) Class0/I MYSOs. 5. From the evolved state of the Sh2-90 \hii region, together with the presence of B-type stars and YSOs embedded in a thin shell of dense gas close to the IF of the \hii region, we suggest that the formation of these sources have possibly been triggered by the expansion of the \hii region. However, detailed velocity and age measurements of the stars in the \hii region could give more insights into this scenario. In summary, it appears that multi-generation star formation is going on, but it remains unclear how the star formation sites and processes are interlinked. Taking the present observational evidence at face value, we suggest that triggered star formation possibly takes place at the immediate periphery of Sh2-90. However, we hypothesize that the MYSOs currently observed at various locations of the complex could have formed spontaneously or by some other processes.	14	3	1403.7925	<BR /> Aims: We investigate the star formation activity in the molecular complex associated with the Galactic H ii region Sh2-90. <BR /> Methods: We obtain the distribution of the ionized and cold neutral gas using radio-continuum and Herschel observations. We use near-infrared and Spitzer data to investigate the stellar content of the complex. We discuss the evolutionary status of embedded massive young stellar objects (MYSOs) using their spectral energy distribution. <BR /> Results: The Sh2-90 region presents a bubble morphology in the mid-infrared. Radio observations suggest it is an evolved H ii region with an electron density ~144 cm<SUP>-3</SUP>, emission measure ~ 6.7 × 10<SUP>4</SUP> cm<SUP>-6</SUP> pc and an ionized mass ~55 M<SUB>⊙</SUB>. From Herschel and CO (J = 3 - 2) observations we found that the H ii region is part of an elongated extended molecular cloud (H<SUB>2</SUB> column density ≥ 3 × 10<SUP>21</SUP> cm<SUP>-2</SUP> and dust temperature 18-27 K) of total mass ≥ 1 × 10<SUP>4</SUP> M<SUB>⊙</SUB>. We identify the ionizing cluster of Sh2-90, the main exciting star being an O8-O9 V star. Five cold dust clumps, four mid-IR blobs around B stars, and a compact H ii region are found at the edge of the bubble. The velocity information derived from CO data cubes suggest that most of them are associated with the Sh2-90 region. One hundred and twenty-nine low mass (≤3 M<SUB>⊙</SUB>) YSOs have been identified, and they are found to be distributed mostly in the regions of high column density. Four candidate Class 0/I MYSOs have been found. We suggest that multi-generation star formation is present in the complex. From evidence of interaction, time scales involved, and evolutionary status of stellar/protostellar sources, we argue that the star formation at the edges of Sh2-90 might have been triggered. However, several young sources in this complex are probably formed by some other processes. <P />Full Table 5 is only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/566/A122">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/566/A122</A>	false	[ "B stars", "multi-generation star formation", "high column density", "total mass ≥", "the Galactic H ii region", "a compact H ii region", "an evolved H ii region", "the H ii region", "emission measure", "cm", "embedded massive young stellar objects", "CO data cubes", "the star formation activity", "Radio observations", "≥", "Sh2", "the main exciting star", "the star formation", "-", "anonymous ftp" ]	10.84814	10.285969	183
482937	[ "Straniero, O.", "Cristallo, S.", "Piersanti, L." ]	2014ApJ...785...77S	[ "Heavy Elements in Globular Clusters: The Role of Asymptotic Giant Branch Stars" ]	65	[ "INAF-Osservatorio Astronomico di Collurania, I-64100 Teramo, Italy; INFN-sezione di Napoli, I-80126 Napoli, Italy;", "INAF-Osservatorio Astronomico di Collurania, I-64100 Teramo, Italy; INFN-sezione di Napoli, I-80126 Napoli, Italy", "INAF-Osservatorio Astronomico di Collurania, I-64100 Teramo, Italy; INFN-sezione di Napoli, I-80126 Napoli, Italy" ]	[ "2014A&A...570A..46C", "2014A&A...572A.108T", "2014ApJ...795...34S", "2014ApJ...797...44F", "2014MNRAS.445.3239P", "2014PASA...31...30K", "2014PhRvL.113s1302A", "2015A&A...578A..33L", "2015AJ....150...63J", "2015ASPC..497..301C", "2015ApJ...801...53C", "2015ApJ...803...12L", "2015ApJS..219...40C", "2015EAS....71..251G", "2015MNRAS.449..506B", "2015MNRAS.450..815M", "2015MNRAS.452.2804S", "2015arXiv151100330F", "2016ApJ...825...26K", "2016ApJ...825...38H", "2016ApJ...833...81H", "2016ApJS..225...24P", "2016EPJA...52...76A", "2016JCAP...05..057G", "2016JPhCS.703a2005E", "2016MNRAS.455.2417R", "2016MNRAS.455.3848J", "2016MNRAS.458.2122D", "2016MmSAI..87..229K", "2016MmSAI..87..289M", "2016PhRvC..93e5803T", "2017A&A...598A.128S", "2017A&A...599A..39A", "2017A&A...601A..96L", "2017ApJ...835...97B", "2017ApJ...846...23O", "2017IAUS..316..267M", "2017MNRAS.465.4817S", "2017MNRAS.471..824B", "2017PhRvC..95b5803U", "2017hsn..book..461K", "2018ApJ...859..105C", "2018EPJWC.18401004C", "2018MNRAS.473..984S", "2018MNRAS.480..538R", "2019MNRAS.490.2219N", "2020A&A...642A.227A", "2020IAUS..351..309K", "2021A&A...645A..10G", "2021ApJ...912...72M", "2021MNRAS.508.3499V", "2021Univ....7..200A", "2022MNRAS.509.4430K", "2022MNRAS.516.3515M", "2022NuPhA102322450N", "2022Univ....8..243M", "2023MNRAS.520.5938S", "2024A&A...683A.138C", "2024ApJS..270...28R", "2024IAUS..377...98M", "2024MNRAS.527.7940M", "2024PASA...41...25M", "2024PhRvC.109b5808S", "2024PhRvL.132l2701A", "2024arXiv240100545C" ]	[ "astronomy" ]	14	[ "Galaxy: abundances", "globular clusters: general", "stars: AGB and post-AGB", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1957PASP...69..201C", "1957RvMP...29..547B", "1958ZA.....46..108B", "1965ApJS...11..121S", "1967ApJ...148...69C", "1968pss..book.....C", "1979ApJ...232..831B", "1980ApJ...237..130S", "1981ApJ...245L..79C", "1983ARA&A..21..271I", "1983MmSAI..54..173D", "1985A&A...150...33B", "1985ApJ...296..204C", "1986ApJ...301..587Y", "1987A&A...182..243M", "1988A&AS...76..157S", "1988ADNDT..40..283C", "1988ApJ...328..653B", "1989ApJS...71...47C", "1989RPPh...52..945K", "1990ApJS...74..463C", "1991ApJ...367..228R", "1991ApJ...370..295C", "1991ApJ...371..665R", "1992A&A...261..164J", "1993ApJ...413..641V", "1993PhRvC..48.2746D", "1994ApJ...437..396K", "1994MNRAS.267..711W", "1995A&A...297..727B", "1995A&A...302..382M", "1995ApJ...440L..85S", "1995ApJ...445L..39C", "1995ApJ...450..302L", "1996A&A...313..497F", "1996ApJ...456..902R", "1996ApJ...464..943I", "1996ApJ...473..383F", "1997A&A...324L..81H", "1997ApJ...478..332S", "1998A&A...340..160C", "1998ApJ...497..388G", "1998ApJ...502..737C", "1998ApJ...508L.103C", "1999A&A...344..617M", "1999ARA&A..37..239B", "1999ASPC..173..133C", "1999ApJ...513L..49F", "1999NuPhA.656....3A", "2000A&A...360..952H", "2000A&A...362..599G", "2000AJ....119.1239S", "2000MmSAI..71..719S", "2001A&A...368..969S", "2001AJ....122.1438I", "2001ApJ...549..346T", "2001ApJ...554.1159C", "2001ApJ...559.1117A", "2001MNRAS.322..231K", "2001PhRvL..87t2501J", "2002A&A...387..507M", "2002A&A...395...69D", "2002ApJ...569..549N", "2002ApJ...581..585P", "2003A&A...409..715W", "2003ApJ...591.1220L", "2003ApJ...592L..67A", "2003ApJ...595.1114Y", "2003MNRAS.342...86W", "2004A&A...421L..25G", "2004ApJ...615..934L", "2005A&A...431..279V", "2005A&A...438..273V", "2005ApJ...631..868D", "2006A&A...457..995V", "2006ApJ...642..225F", "2006NuPhA.777..311S", "2007A&A...462.1051P", "2007ApJ...658.1203J", "2007ApJ...667..489C", "2007ApJ...667L..81S", "2007M&PS...42.1103B", "2007MmSAI..78..543P", "2008ARA&A..46..241S", "2008ApJ...687L..95P", "2008ApJ...689.1031Y", "2009A&A...506.1277G", "2009A&A...508.1539M", "2009ApJ...694..971A", "2009ApJ...696..797C", "2010ApJ...715L..94A", "2010ApJ...722.1373J", "2010MNRAS.403.1413K", "2011A&A...534A..29D", "2011ApJ...729...40L", "2011ApJ...737L...8A", "2011ApJ...742...37R", "2011ApJS..197...17C", "2011PhRvC..83e2802B", "2012A&A...538L...2F", "2012A&A...540A..44V", "2012A&A...541A..15M", "2012A&A...547C...2G", "2012A&ARv..20...50G", "2012ApJ...747....2L", "2012MNRAS.422..849B", "2013A&A...560A..74D", "2013ApJ...763...22D", "2013ApJ...774...98P", "2013ApJ...776...59D", "2013MNRAS.433..366D", "2013MNRAS.433.1941L", "2013PhRvC..87d5805B", "2013PhRvL.111k2501R", "2014ApJ...780...95L", "2014MNRAS.437..195D" ]	[ "10.1088/0004-637X/785/1/77", "10.48550/arXiv.1403.0819" ]	1403	1403.0819_arXiv.txt	All the elements heavier than iron are mainly produced by neutron captures\footnote {A few isotopes are actually synthesized by the so-called p process whose overall contribution to the elemental abundances is, however, rather small.}. There exist two different nucleosynthesis processes of this type, the slow (s) process and the rapid (r) process \citep{b2hf}. Since the typical neutron density of the r process is more than 10 orders of magnitude larger than that of the s process, significantly different physical conditions are implied and, in turn, very different astrophysical environments. The r process is commonly associated with massive stars. Two are the proposed scenarios: core-collapse supernovae (type II, Ib and Ic) and Neutron-Star mergers. Although none of the proposed astrophysical sites has been confirmed by direct observations, the yields of the r process are commonly found in all the Galactic components, very-metal-poor stars included. Such a prompt pollution, demonstrates that the r process takes place in stars that evolves on a very short timescale \citep[see][]{sneden2008}. On the contrary, our knowledge of the s-process site has been greatly improved in the last 20 years \citep[for a review see][]{busso1999}. First of all, it should be reminded that the most abundant products of the s process are the so-called {\it neutron-magic nuclei}, whose neutron-capture cross section is particularly low compared to the cross section of nearby nuclei. When the s-process flow encounters a neutron-magic nucleus, it acts as a bottleneck, so that its abundance is greatly enhanced with respect to the nearby non-magic nuclei, for which a nearly local equilibrium is established, as given by: $\sigma_A N_A=\sigma_{A-1} N_{A-1}$ \footnote{ $\sigma_A=\frac{< \sigma v>}{v_{th}}=\frac{2}{\sqrt{KT}}\int_{0}^{\infty}E\sigma_{n}(E)exp\left ( -\frac{E}{KT} \right )dE$ is the Maxwellian averaged cross section (MACS) and $N_A$ is the fraction of isotopes with atomic mass A.}. The most important neutron-magic nuclei encountered by the s process are $^{88}$Sr, $^{89}$Y, $^{90}$Zr, $^{138}$Ba, $^{139}$La, $^{140}$Ce, $^{141}$Pr, $^{142}$Nd, $^{208}$Pb and $^{209}$Bi. Each of these nuclei corresponds to a peak in the distribution of the solar system abundances. The first three are the major contributors to the light-s peak, while those from $^{138}$Ba to $^{142}$Nd contribute to the heavy-s peak. As pointed out in the seminal paper of \citet{b2hf}, the s process follows simple general rules. Three are the main players: neutrons (or neutron sources), seeds (Fe nuclei) and neutron poisons. The latter are light elements that compete with the seeds in the neutron-capture nucleosynthesis. In this context, a fundamental quantity that characterizes the s process is the neutron-to-seed ratio, i.e., $f=\frac{neutrons-poisons}{seeds}$, where {\it neutrons, poisons and seeds} represent fractions by number. As firstly shown by \citet{cameron1957}, the synthesis of the heaviest elements, such as Pb, requires a relatively large value of this ratio ($f>20$), while for low values, namely $f\sim 1$, only light-s are produced. Note that the number of seeds directly scales with the metallicity, so that the production of the heaviest s elements is generally favored at low Z \citep{busso1999,cristallo2009}. As a matter of fact, the cosmic concentration of lead is the result of the pollution caused by low-metallicity AGB stars \citep[e.g.,][]{travaglio2001}. Other important quantities that characterize the s process are the neutron density ($N_n$), the temperature (T) and the timescale (i.e., the duration of the s-process episode). They determine the time integrated neutron flux, or neutron exposure, namely $\tau = \int N_nv_{th}dt$\footnote{$v_{th}$ is the thermal velocity, which depends on T}. Note that the larger the neutron exposure the larger the probability to overshoot the neutron-magic nuclei. Moreover, the Maxwellian averaged cross sections depend on the temperature, while the neutron density is important for the various branchings occurring along the s-process path. Indeed, when a neutron capture produces an unstable nucleus, the $\beta$ decay may compete with a further neutron capture \citep[see][]{kappeler1989}. For each branching, it exists a critical value of the neutron density given by the ratio of the decay rate and the neutron-capture rate. When the neutron density is much larger than this critical value, the neutron capture is favored with respect to the $\beta$ decay, while the opposite occurs at low neutron density. In this way, the neutron density determines the abundances of the isotopes on the alternative paths. Some examples are the branchings at $^{79}$Se, $^{85}$Kr, $^{95}$Zr, $^{134}$Cs and $^{151}$Sm. In general, those isotopes/elements whose production is sensitive to the neutron density are good estimators of the physical conditions of the s-process site \citep{lambert1995,abia2001,aoky2003,barzyk2007,vanraai2012,lugaro2014}. By analyzing the heavy element composition of the solar system, three different components of the s process have been formerly identified, namely the weak, the main and the strong \citep{seeger1965,clayton1967}. Each s-process component implies a specific range of neutron exposures and, in turn, a specific range of the quantities characterizing different s-process sites, i.e., $f$, $N_n$, T and the timescale. The weak component, which includes nuclei with $29<Z<40$, is synthesized in the He-burning core and, later on, in the C-burning shell of massive stars \citep[M$> 10$ M$_\odot$,][]{raiteri1991, raiteri1991b, kappeler1994, pignatari2006}. The neutron density may vary from $\sim 10^6$ neutrons/cm$^3$, in the case of the He burning, up to $\sim 10^{11}$, for the C-burning. Temperatures and timescales are also very different, but the neutron exposure is similar, namely $\sim0.06$ mbarn$^{-1}$. In both cases, neutrons are provided by the \nean reaction, so that $^{22}$Ne is a necessary ingredient for the weak-s process. In practice, $^{22}$Ne is synthesized during He burning, through the $^{14}$N$(\alpha,\gamma)^{18}$F$(\beta^+)^{18}$O$(\alpha,\gamma)^{22}$Ne chain, where $^{14}$N is that left by the former CNO burning. Therefore, the fraction (by number) of $^{22}$Ne nuclei available for the s process in massive stars is approximately equal to the original fraction of C+N+O nuclei. Such an occurrence implies that the synthesis of the weak component is less efficient at low Z, because of the paucity of C+N+O and, in turn, of $^{22}$Ne. For instance, during core-He burning, the main neutron poison is $^{25}$Mg that is secondary like, since it is directly produced by the \nean reaction. As a result, the weak process yields decrease roughly linearly with metallicity. Instead, in the C-burning shell there are primary like neutron poisons (e.g., $^{16}$O, $^{23}$Na, $^{24}$Mg), which do not depend on the metallicity. Therefore, the s-process efficiency in the C-burning shell is strongly suppressed at low Z \citep{pignatari2007}. Recently, \citet{pignatari2008} show that in very low-metallicity fast-rotating massive stars, fresh C synthesized by the 3$\alpha$ reaction may be transported by meridional circulation into the H-rich envelope, thus increasing the amount of C+N+O. They find that this phenomenon would allows an efficient s-process nucleosynthesis, up to Pb. However, in a more recent paper, \citet{frisch2012} argue that this result is due to the use of a particularly low rate of the $^{17}$O($\alpha,\gamma)^{21}$Ne reaction, i.e., that suggested by \citet{desc1993}, which is up to a factor of 1000 lower than the values reported in the widely used reaction rate compilations \citep{CF88, NACRE}. In the He-burning core of fast rotating massive stars, this reaction is expected to destroy most of the $^{17}$O released by the poisoning reaction $^{16}$O(n,$\gamma)^{17}$O. The suppression of the $^{17}$O($\alpha,\gamma)^{21}$Ne favors the competitive channel $^{17}$O($\alpha,$n)$^{20}$Ne, so that the neutrons subtracted by the $^{16}$O would be recycled. However, new experiments reinvestigated both channels of the $^{17}$O$+\alpha$, confirming previous findings \citep{best2011, best2013}. In particular, they find that the $\gamma$ channel is strong enough to compete with the neutron channel, thus leading to a less efficient neutron recycling. Fast-rotating massive stars might still play a role in the production of the weak component (up to Sr), but no significant s-process contribution to heavier elements are expected \citep[see Figure 14 in][]{best2013}. The main and the strong components, which include nuclei with $37<Z<84$, are produced by low-mass stars ($1.5<$M/M$_\odot<2.5$) \citep{straniero1995, gallino1998, cristallo2009, cristallo2011}. In these stars, recursive thermonuclear runaways of the shell-He burning, called thermal pulses (TPs), take place during the AGB phase. Two important events are connected to the occurrence of these thermal pulses. First of all, owing to the excess of nuclear energy released by the thermonuclear runaway, an extended convective instability takes place within the He-rich layer. Later on, owing to the expansion powered by He burning, the shell-H burning dies down and the inner border of the convective envelope can attain the He-rich zone (third dredge up - TDU). The s-process nucleosynthesis in low-mass stars mostly occurs during the relatively long interpulse period ($\sim 10^5$ yr), namely the time elapsed between two subsequent thermal pulses, in a thin radiative layer located at the top of the He-rich zone \citep{straniero1995}. This layer is known as the $^{13}$C pocket, because it is enriched in $^{13}$C. The neutron source is the \ctan reaction, which requires a temperature of $\sim90-100$ MK and releases low-density neutron fluxes, i.e., about $10^7$ neutrons/cm$^{3}$, and neutron exposures between 0.1 and 0.4 mbarn$^{-1}$ \citep{gallino1998}. A second neutron burst giving rise to a higher neutron density ($>10^{11}$ neutrons/cm$^{3}$) is due to the marginal activation of the \nean reaction within the convective zone generated by a thermal pulse, where the temperature may exceed 300 MK. In this case, the timescale is rather short ($\sim 1$ yr), so that the resulting neutron exposure is lower than that of the first neutron burst. These low-mass stars are the main contributors to the s-process elements in the solar system. However, because of their long lifetime ($\ge1$ Gyr), it appears that they cannot have contaminated the gas from which the galactic halo formed. This is the standard paradigm for the heavy element composition of the halo. In practice, only r-process yields are expected in fossil records of the early Galaxy, the s process being hampered by the secondary nature of the neutron sources in massive stars (weak component) and by the too long lifetime of low-mass AGBs (main and strong components). Spectroscopic studies generally confirm such a scenario: single halo stars are r-process enriched, but s-process poor \citep[see][and references therein]{sneden2008}. Exceptions are the CEMP-s (Carbon-Enhanced-Metal-Poor stars, where the ``s'' stay for s-rich). In this case however, the s and the C enrichments are a consequence of mass transfer or wind accretion in binary systems, a process occurring on a longer timescale \citep[see][and references therein]{bisterzo2012, lugaro2012}. In this context, recent spectroscopic studies of Globular Clusters (GCs) revealed a rather different scenario. While the r-process yields generally appear similar to those observed in halo field stars, some GCs show a clear signature of the s-process main component pollution. The few GC stellar populations where an s-process enrichment has been discovered are: M4 \citep{yong2008, dorazi2013a}, $\omega$-Cen, only stars with [Fe/H]$>$-1.6 \citep{smith2000, johnson2010, dorazi2011}, and the redder main sequences of M22 \citep{roederer2011} and NGC1851 \citep{gratton1851}. Recently, s-process overabundances have been also found in M2 stars \citep{lardo2013}. Other clusters, like M5 \citep{ivans2001, yong2008}, as well as the most metal-poor stellar populations of $\omega$-$Cen$, M22 and NGC1851, present a ``normal'' halo distribution of the heavy elements characterized by a pure r-process pollution. These challenging observations represent a further evidence of the existence of multiple stellar populations in GCs. Nevertheless, at variance with other spectroscopic anomalies, such as the O-Na anticorrelation \citep[][and references therein]{gratton2012}, the s-process enhancement is not a common feature of the majority of the GCs in the Milky Way. Therefore, a different class of polluters should be responsible for the heavy-element anomalies. Such a conclusion is also supported by the fact that in M22 and NGC1851 the O-Na anticorrelation is observed in both s-rich and s-poor stars of the same cluster. Moreover, all stars in M4 show a similar overabundance of the s elements but this enrichment is uncorrelated with the [Na/Fe]. More intriguing, some spectroscopic indexes, which depend on the metallicity of the polluters, do not match the theoretical expectations for low-mass AGB stars, which are considered the most important producers of the galactic s-process main and strong components. In particular, the ratio between heavy-s (Ba, La or Nd) and light-s elements (Sr, Y or Zr) are found in solar proportions ([hs/ls]$\sim0$]), while an excess of heavy-s is expected at low Z. Therefore, the polluters responsible for such a peculiar chemical pattern cannot be the same stars responsible for the bulk of the s-process yields in the Galaxy. In this paper we study the characteristics of the s-process nucleosynthesis in metal-poor AGB stars of low and intermediate mass. We will discuss, in particular, the variations of the nucleosynthesis outcomes with the stellar mass. In the next section we review the most important inputs physics and how they are included in our stellar evolution code. In section 3 we analyse the operation of the two neutron sources active in thermally pulsing AGB stars. This analysis is based on the models presented in section 4. The theoretical yields we derive from these models may be used to test various scenarios for GC formation that have been proposed to explain photometric and spectroscopic evidences of multiple populations, among which: multiple photometric sequences, star-to-star variations of the chemical composition, which cannot be ascribed to internal physical processes, and anomalous color dispersion of horizontal branch stars or the so-called second parameter problem \citep[for a recent review see][]{gratton2012}. Several hypotheses about the GC formation have been proposed to explain the new observational framework, such as: inhomogeneities of the primordial material, merging of smaller stellar systems, pollution with external material felt into the gravitational potential well of the cluster and various self-pollution scenarios. Which of these scenarios can also provide an explanation for the s-process enhancements observed in a few GC stellar populations? Which stars are responsible for the s-process contamination in GCs? What are the special conditions determining the onset of this peculiarity? These issues are addressed in section 5 and 6. We show, in particular, that AGB stars with mass ranging between 3 to 6 M$_\odot$ can produce the yields necessary to reproduce the observed heavy-element anomalies. In this case, we find that the time elapsed between the formation of the polluters and that of the polluted stellar populations should be of the order of 100-200 Myr.	In this paper we presented new models of low-metallicity AGB stars with mass in the range 2-6 M$_\odot$. The heavy elements yields of these models allow us to reproduce most of the observed features of the s-process main and strong components, as shown by stars of some GC stellar populations. The comparison between the theoretical predictions and the observed overabundance of s elements has been done by adopting a simple MP model for the early GC history. This model implies two main temporal steps, namely: 1) a first stellar generation forms from pristine gas whose heavy element composition is that typical of the bulk of the galactic halo, i.e., r rich, but s poor. 2) after about $150\pm 50$ Myr, a second stellar generation forms within a newborn GC from the gas ejected by the stars of the first generation, possibly diluted with some amount of pristine gas. The first generation may or may not be a member of the cluster where the second generation is observed. In other words, the pollution may be either primordial or internal to the cluster (self pollution). According to this picture, if the star formation definitely halts in less than $\sim 50$ Myr, namely before that intermediate-mass stars (M$\le6$ M$_\odot$) have time to evolve up to the AGB phase and pollute the interstellar gas, the GC will be s-process poor. This occurrence explains why s-process enhancements are so rare in GCs. It also implies that the more massive stars, whose lifetime is shorter than 50 Myr, do not substantially contribute to the main and strong components of the s process. On the contrary, these stars should be responsible, fully or partially, for the more common variations of C, N, O, Na, Mg, Al and other light elements. For this reason, a more powerful and complete pollution model may be obtained by coupling the yields here presented to those of more massive stars. On the other hand, physical phenomena not yet included in the present stellar models, such as rotation, may also improve the theoretical tool. We are grateful to F. Kappeler and I. Dillman, for they help in interpreting the KADONIS reaction rates, and to D. Yong for providing us the M4 and M5 data in a computer readable form. The present work has been support by the PRIN-INAF 2010 and FIRB-MIUR 2008 (RBFR08549F-002) programs. Extended Tables of the models presented in this paper are available in the FRUITY database (fruity.oa-teramo.inaf.it).	14	3	1403.0819	Recent observations of heavy elements in globular clusters reveal intriguing deviations from the standard paradigm of the early galactic nucleosynthesis. If the r-process contamination is a common feature of halo stars, s-process enhancements are found in a few globular clusters only. We show that the combined pollution of asymptotic giant branch (AGB) stars with a mass ranging between 3 to 6 M <SUB>⊙</SUB> may account for most of the features of the s-process overabundance in M4 and M22. In these stars, the s process is a mixture of two very different neutron-capture nucleosynthesis episodes. The first is due to the <SUP>13</SUP>C(α, n)<SUP>16</SUP>O reaction and takes place during the interpulse periods. The second is due to the <SUP>22</SUP>Ne(α, n)<SUP>25</SUP>Mg reaction and takes place in the convective zones generated by thermal pulses. The production of the heaviest s elements (from Ba to Pb) requires the first neutron burst, while the second produces large overabundances of light s (Rb, Sr, Y, Zr). The first mainly operates in the less massive AGB stars, while the second dominates in the more massive. From the heavy-s/light-s ratio, we derive that the pollution phase should last for 150 ± 50 Myr, a period short enough compared to the formation timescale of the globular cluster system, but long enough to explain why the s-process pollution is observed in a few cases only. With few exceptions, our theoretical prediction provides a reasonable reproduction of the observed s-process abundances, from Sr to Hf. However, Ce is probably underproduced by our models, while Rb and Pb are overproduced. Possible solutions are discussed.	false	[ "globular clusters", "light s", "the observed s-process abundances", "the s-process overabundance", "thermal pulses", "intriguing deviations", "large overabundances", "Zr", "few exceptions", "Sr", "the early galactic nucleosynthesis", "asymptotic giant branch", "the s process", "the globular cluster system", "second", "M22", "halo stars, s-process enhancements", "a few globular clusters", "Hf", "the s-process pollution" ]	8.518925	9.05756	-1
743560	[ "Böhringer, H." ]	2014MmSAI..85..396B	[ "X-ray observations of the chemical abundances in the Intra-Cluster Medium" ]	1	[ "Max-Planck-Institut für extraterrestrische Physik Giessenbachstrasse, 85748 Garching Germany," ]	[ "2017ApJ...843..105S" ]	[ "astronomy" ]	2	[ "Galaxies: clusters: intracluster medium", "Galaxy: abundances", "X-rays: galaxies: clusters", "Astrophysics - Cosmology and Extragalactic Astrophysics" ]	[ "1984ApJ...286..644N", "1986RvMP...58....1S", "1989A&A...215..147B", "1989GeCoA..53..197A", "1995MNRAS.275..720N", "1999ApJS..125..439I", "1999MNRAS.302...22C", "2000ApJ...528..139D", "2001A&A...365L.181B", "2001ApJ...551..153D", "2001ApJ...556L..91S", "2002A&A...381...21F", "2003A&A...401..443M", "2003ApJ...591.1220L", "2003ApJ...595..151B", "2003ApJ...598..250S", "2004A&A...416L..21B", "2004A&A...419....7D", "2004A&A...420..135T", "2004A&A...420..833X", "2005AdSpR..36..677B", "2006A&A...459..353W", "2006MNRAS.371.1483S", "2006MNRAS.372.1840R", "2007A&A...462..429B", "2007A&A...462..953M", "2007A&A...465..345D", "2008ApJS..174..117M", "2009ARA&A..47..481A", "2009ApJ...705L..62T", "2010A&ARv..18..127B", "2011ARA&A..49..409A", "2012A&A...537A.142B", "2013AN....334..416D" ]	[ "10.48550/arXiv.1403.5261" ]	1403	1403.5261_arXiv.txt	Clusters of galaxies are the largest clearly defined objects in our Universe. They comprise masses in the range of about $10^{14}$ to $ 3 \times 10^{15}$ M$_{\odot}$. In the mass range of about few $10^{13}$ to $10^{14}$ M$_{\odot}$ we find the groups and poor clusters of galaxies. As an integral part of the cosmic large-scale structure clusters form from large-scale overdense regions in the matter distribution, the seeds of which have been set presumably at the epoch of inflation. Galaxy clusters form very late in the history of our Universe and the bulk of the cluster population has emerged only after a redshift of 2. Before, hardly any massive clusters existed but groups of galaxies have been present \citep{boe10,all12} \begin{figure}[t!] \resizebox{\hsize}{!}{\includegraphics[clip=true]{Boehringer_fig1.pdf}} \caption{\footnotesize XMM-Newton image and spectrum of the central region of the X-ray halo of M87 in the Virgo cluster. The image shows the X-ray surface brightness of a region with a size of about 70 kpc radius. Overlayed is the X-ray spectrum of the inner 20 kpc radius region showing lines of the most abundant heavy elements O, Mg, Si, S, Ar, Ca, Fe, and Ni. The X-ray emission represented by the spectrum comes from the entire volume of the X-ray emitting region seen in the image in projection. } \label{fig1} \end{figure} Most of the mass of galaxy clusters, about 84\% in massive systems, is made up by the so-called Dark Matter, whose nature we still don't know. Only about 4\% is made up by the stars in galaxies and 12\% by the hot intracluster plasma seen in X-rays. The formation of a galaxy cluster is thus mostly determined by the gravitational dynamics of the dark matter, which forms through gravitational collapse a virialised, nearly spherical symmetric system, which can approximately be described by e.g. a NFW model \citep{nfw95}. Galaxies which mostly have been formed before the cluster collapse and the intergalactic medium are collapsing simultaneously with the Dark Matter. While the galaxy population gains a velocity dispersion of the order of 1000 km s$^{-1}$ in massive systems, the gas heats up to temperatures of several ten Million degrees and forms the intracluster medium (ICM). The ICM emits thermal radiation in the form of soft X-rays, in just the wavelength regime where X-ray telescopes show their best performance. Thus galaxy clusters are very rewarding objects for imaging X-ray astronomy \citep{sar86, boe10}. This fact, that the cluster ICM can be observed with modern X-ray telescopes, makes galaxy clusters unique astrophysical laboratories, where we can study all the ``baryonic matter'', in the galaxies and in the ICM simultaneously in detail. In the extracluster space, the intergalactic medium is mostly invisible and we can study it essentially only through absorption effects. A most important finding comparing ICM and galaxies in clusters is the fact that in massive systems we find more mass in heavy elements (metals) in the ICM than in the galaxies. Thus the galaxy cluster laboratories provide us with the unique opportunity to take a full account of all metals produced by stellar nucleosynthesis in a certain, representative volume of the Universe. \begin{figure}[t!] \resizebox{\hsize}{!}{\includegraphics[clip=true]{Boehringer_fig2.pdf}} \caption{\footnotesize Theoretically calculated X-ray spectra of hot plasma with solar element abundances at temperatures of $10^7$ K (left) and $10^8$ K (right). The blue, green, and red lines give the contribution of Bremsstrahlung, recombination continuum, and two-photon emission. Some major lines are labled by the element responsible for the line emission. The red bar on top of the figure shows the approximate spectral range covered the XMM-Newton or Chandra detectors \citep{boe89, bow10}. } \label{fig2} \end{figure} The X-ray emitting ICM which is heated by shocks and compression during cluster formation keeps its heat even over a Hubble time, apart from a small central region in cool core clusters where cooling is more efficient. The observed X-ray emission is therefore that of hot plasma effectively in thermal equilibrium and so far no deviation from thermal equilibrium has been detected. Therefore the X-ray spectrum is readily modeled as explained in the next section. X-ray observatories like ESA's XMM-Newton and NASA's Chandra space observatories detect X-ray photons as single events taking a record of the direction and energy of incidence. Therefore the photons can be binned in the form of images as well as into spectra. Fig. 1 shows as an example the X-ray image and spectrum of the central region of the X-ray halo of the giant elliptical galaxy M87 in the center of the Virgo cluster \citep{boe01}.	The heavy element abundances in the ICM of clusters of galaxies provide a lot of interesting and important information on the history of the nucleosysthesis of these metals, on their sources in the clusters and on the possible transport processes. They also allow us to test the supernova models through their predicted abundance yields. But better observational data are needed. Significant progress can be expected, when the new X-ray calorimeter detectors become available with much better spectral resolution and well resolved lines. The first space mission that can provide this capability is ASTRO-H which will hopefully be launch successfully soon.	14	3	1403.5261	Clusters of galaxies as the largest clearly defined objects in our Universe are ideal laboratories to study the distribution of the most abundant chemical elements heavier than hydrogen and helium and the history of their production. The cluster environment allows us to study the element abundances not only inside the galaxies, but also in the intergalactic space, the intracluster medium. Since the intracluster medium is heated to temperatures of several ten Million degrees, we can study the chemical composition of this medium through X-ray spectroscopy. Up to 13 heavy elements have been detected by X-ray spectroscopy so far. The element most easily detected in the X-ray spectra is iron. In massive galaxy clusters we find a larger mass of heavy elements in the intracluster medium than in the galaxies. The consideration of the intracluster medium is therefore vital for an understanding of the complete history of nucleosynthesis of the heavy elements. The observed abundances for all elements heavier than nitrogen can roughly be modeled by using two types of sources: core collapse supernovae and supernovae type Ia. So called cool-core galaxy clusters show a larger heavy element abundance in the cluster center which seems to be enriched primarily by products of supernovae of type Ia. The evidence for observations of an evolution of the heavy element abundance with redshift has still a moderate significance.	false	[ "heavy elements", "type Ia.", "X-ray spectroscopy", "massive galaxy clusters", "core collapse", "a larger heavy element abundance", "type", "galaxies", "the heavy element abundance", "cool-core galaxy clusters", "Ia.", "Clusters", "the heavy elements", "ideal laboratories", "the element abundances", "iron", "nucleosynthesis", "sources", "the intracluster medium", "The element" ]	13.376619	5.431157	134
483107	[ "Rathborne, J. M.", "Longmore, S. N.", "Jackson, J. M.", "Foster, J. B.", "Contreras, Y.", "Garay, G.", "Testi, L.", "Alves, J. F.", "Bally, J.", "Bastian, N.", "Kruijssen, J. M. D.", "Bressert, E." ]	2014ApJ...786..140R	[ "G0.253+0.016: A Centrally Condensed, High-mass Protocluster" ]	68	[ "CSIRO Astronomy and Space Science, P.O. Box 76, Epping NSW 1710, Australia", "European Southern Observatory, Karl-Schwarzschild-Str. 2, D-85748 Garching bei Munchen, Germany; Current address: Astrophysics Research Institute, Liverpool John Moores University, Egerton Wharf, Birkenhead CH41 1LD, UK.", "Institute for Astrophysical Research, Boston University, Boston, MA 02215, USA", "Institute for Astrophysical Research, Boston University, Boston, MA 02215, USA; Current address: Department of Astronomy, Yale University, P.O. Box 208101, New Haven, CT 06520-8101, USA.", "CSIRO Astronomy and Space Science, P.O. Box 76, Epping NSW 1710, Australia", "Universidad de Chile, Camino El Observatorio1515, Las Condes, Santiago, Chile", "European Southern Observatory, Karl-Schwarzschild-Str. 2, D-85748 Garching bei Munchen, Germany; INAF-Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy; Excellence Cluster Universe, Boltzmannstr. 2, D-85748 Garching, Germany", "University of Vienna, Türkenschanzstrasse 17, A-1180 Vienna, Austria", "Center for Astrophysics and Space Astronomy, University of Colorado, UCB 389, Boulder, CO 8030, USA", "Astrophysics Research Institute, Liverpool John Moores University, Egerton Wharf, Birkenhead CH41 1LD, UK", "Max-Planck Institut fur Astrophysik, Karl-Schwarzschild-Strasse 1, D-85748 Garching, Germany", "CSIRO Astronomy and Space Science, P.O. Box 76, Epping NSW 1710, Australia" ]	[ "2014A&A...568A..56J", "2014ApJ...792L..14J", "2014ApJ...795...28B", "2014ApJ...795L..25R", "2014CQGra..31x4006K", "2014MNRAS.440.3370K", "2014prpl.conf..291L", "2015ApJ...802..125R", "2015ApJ...802L..15Z", "2015ApJ...805...72M", "2015ApJ...806..130T", "2015ApJ...815..130G", "2015ApJS..219....2J", "2015EAS....75...43L", "2015MNRAS.447.1059K", "2015MNRAS.448.1847H", "2015MNRAS.449..715W", "2015MNRAS.451.3679B", "2015MNRAS.453..739K", "2016ApJ...832..143F", "2016ApJ...833..248Y", "2016MNRAS.455.3763B", "2016MNRAS.457.2675H", "2016MNRAS.457.4536W", "2016MNRAS.461L..16M", "2016PASA...33...30R", "2017A&A...599A.136L", "2017A&A...603A..89K", "2017ApJ...839....1L", "2017ApJ...850...77K", "2017IAUS..322...64K", "2017IAUS..322...85H", "2017IAUS..322..147B", "2017MNRAS.466..340C", "2017MNRAS.466.1213K", "2017MNRAS.469.2263B", "2017MNRAS.470.1462L", "2017MNRAS.472.4750D", "2017arXiv170902539R", "2018ApJ...853..171G", "2018ApJ...859...86T", "2018ApJ...860L..14K", "2018ApJ...869..102J", "2018ApJS..236...40T", "2018Galax...6...55L", "2018MNRAS.473.2899P", "2018MNRAS.478.3380J", "2019A&A...628A..27B", "2019ApJ...878..146S", "2019MNRAS.484.5734K", "2019MNRAS.485.1775I", "2019MNRAS.485.2457H", "2019MNRAS.486..283B", "2019MNRAS.486.3307D", "2020ApJ...903..111T", "2020ApJS..251...14H", "2020MNRAS.494..624K", "2020NatAs...4.1064H", "2021A&A...651A..88N", "2021ApJ...908L..31O", "2021ApJ...922...79H", "2021MNRAS.503...77W", "2021NewAR..9301630B", "2022MNRAS.509.4758H", "2023ASPC..534...83H", "2023ApJ...959...36G", "2023MNRAS.520.2245P", "2024ApJ...962...14H" ]	[ "astronomy" ]	11	[ "dust", "extinction", "infrared: ISM", "ISM: clouds", "radio lines: ISM", "stars: formation", "Astrophysics - Astrophysics of Galaxies" ]	[ "1981A&A...103..197G", "1982A&A...115..185W", "1986A&A...155..129W", "1986ApJ...307..350B", "1992A&A...256..595S", "1994ApJ...424..189L", "1995ASPC...77..433S", "1997ApJ...480..235E", "1998ApJ...507..794L", "1999ApJ...525..750F", "1999PASP..111.1049G", "1999RPPh...62..143W", "2000prpl.conf..299K", "2001A&A...370L..49M", "2001ApJ...548L..65M", "2001ApJ...550..761L", "2002ARA&A..40...27C", "2002ApJ...578..211S", "2002MNRAS.337L..17R", "2005A&A...432..369S", "2005A&A...434..949C", "2005JMoSt.742..215M", "2005PASA...22...62L", "2006A&A...454L..13G", "2006ApJ...643.1166F", "2007A&A...468..627V", "2007ApJ...655L..45M", "2008MNRAS.390..759B", "2009A&A...504..415S", "2009ApJ...700L..99A", "2009MNRAS.395.1021L", "2010A&A...516A..78N", "2010ARA&A..48..431P", "2011A&A...525A.116G", "2011A&A...527A..88V", "2011ApJ...735L..33M", "2011ApJS..197...25F", "2011MNRAS.411.1386D", "2011MNRAS.416..178B", "2012A&A...548A.120I", "2012ApJ...746..117L", "2012ApJ...750L..44K", "2012ApJ...758L..29G", "2012MNRAS.419.2961J", "2012MNRAS.420..340K", "2012MNRAS.426.3008K", "2013A&A...549A..45C", "2013A&A...550A.135A", "2013ApJ...765L..35K", "2013ApJ...767L..13R", "2013ApJ...768L..34C", "2013JPCA..117.9404Y", "2013MNRAS.429..987L", "2013MNRAS.433L..15L", "2013MNRAS.435.2598K", "2013PASA...30...57J", "2014MNRAS.440.3370K" ]	[ "10.1088/0004-637X/786/2/140", "10.48550/arXiv.1403.0996" ]	1403	1403.0996_arXiv.txt	\cloud, an infrared dark cloud (IRDC) located in close proximity to the Galactic Centre \citep{Lis94,Lis98,Lis01,Immer12}, has recently gained attention as a potential candidate of a high-mass cluster in a very early stage of formation \citep{Longmore12}. \cloud\, is clearly seen as an extinction feature from mid- to far-IR wavelengths. Recent modelling of \Herschel\, data reveals a central dust temperature of $\sim$20\,K, peak \hh\, column density of $\sim$ 3.3 $\times$ 10$^{23}$\,\cms, and mass of $\sim$ 1.3$\times$ 10$^{5}$\,\Msun\, \citep{Molinari11,Longmore12}. Despite its high mass and density, it shows little evidence for wide-spread star formation, consistent with the low dust temperature. This combination of properties, little star formation and a low dust temperature, yet with a high mass and column density, make \cloud\, extreme when compared to other known Galactic molecular clumps. Indeed, \cite{Ginsburg12} find no other examples of cold, starless clumps with masses $>$ 10$^{4}$\,\Msun\, in the recent Bolocam Galactic Plane Survey. Because it contains $\sim$ 10$^{5}$\,\Msun\, of material, \cloud\, has sufficient mass to form a young massive cluster (YMC) through direct collapse. As such, its detailed study may reveal the initial conditions within a protocluster and the processes by which a high-mass cluster is formed. Young massive clusters are gravitationally bound stellar systems with masses $>$10$^{4}$\,\Msun\,and ages $<$ 100\,Myr \citep{Portegies-Zwart10}. Only a few YMCs have been identified within our Galaxy (e.g., Arches, Quintuplet, Westerlund 1, RSCG, GLIMPSE-CO1: \citealp{Figer99,Clark05,Figer06,Davies11}). Recent work suggests that YMCs may be the `missing link' between open clusters and globular clusters and, as such, may be the local-universe analogs of the progenitors of globular clusters (e.g.\, \citealp{Elmegreen97,Bastian08,Kruijssen12}). Characterising and understanding how these YMCs form is critical to connect Galactic and extra-galactic cluster formation. Identifying a sample of Galactic molecular clouds that may form YMCs is important because their detailed study can shed light not only on how these high-mass clusters form, but potentially on how all clusters form. To date, \cloud\, is one of a handful of known molecular clumps with enough mass to form a cluster of similar mass to these Galactic YMCs through direct collapse. Because \cloud\, shows evidence for structure on small spatial scales (7.5\arcsec), \cite{Longmore12} speculate that it may be destined to form a YMC through hierarchical fragmentation, in a scaled up version of open cluster formation. Indeed, recent models predict that \cloud\, should form a cluster through hierarchical fragmentation in which $>$80\% of the stars should remain bound after the expulsion of the residual gas by feedback \citep{Kruijssen12b}. Located at a distance of $\sim$8.5 kpc, \cloud\, lies $\sim$100\,pc from the Galactic Centre (GC; \citealp{Molinari11}). Its location within the harsh GC environment may provide clues to the formation of such a high-mass protocluster and whether star formation can progress within it. What remains unknown is if the influence of strong tidal forces, high turbulence, and extreme magnetic fields, radiation fields, and cosmic ray fluxes within the GC region helps or hinders the formation of a high-mass cluster. While some YMCs within the Galaxy are located close to the GC (e.g. Arches, Quintuplet), recent work finds that given the large reservoir of dense gas available, the broader GC region appears to be under-producing stars compared to commonly assumed relations between gas mass and star formation (e.g., \citealp{Longmore13}). Molecular line studies show that the bulk of the gas in the Central Molecular Zone (CMZ), the region within $\sim$ 200\,pc of the Galactic Centre, has temperatures of $\sim$ 80\,K and densities $> 10^{4}$\,\cmc\, \citep{Walmsley86, Ao13}. Somewhat more surprising is that many complex organic molecules show bright, widespread emission across the CMZ, leading to the speculation that shock chemistry might dominate in the region \citep{Wilson82,Martin-Pintado01}. Molecular clouds located in the CMZ have exceptional physical properties; they are denser, warmer, more turbulent, and more massive compared to molecular clouds in the Galactic disk. Such high densities are required for the molecular clouds to survive in the steep Galactic Center gravitational potential \citep{Bania86}. Previous molecular line observations toward \cloud\, (also known as GCM0.25+0.011 and M0.25+0.01; \citealp{Gusten81,Lis94}) reveal large line-widths, indicating a high degree of turbulence, similar to other molecular clouds near the GC. Given its location in the CMZ where the interstellar radiation field (ISRF) and cosmic ray ionisation rate (CRIR) are high \citep{Clark13,Yusef-zadeh13}, one might expect a high-density clump like \cloud\, to be externally heated. Indeed, observations have shown that it has a low internal luminosity and that its derived dust temperature increases smoothly from $\sim$ 19\,K in its centre to $\sim$ 27\,K toward its edges \citep{Lis01,Longmore12}. These observations are consistent with recent SPH modelling of the dust and gas temperature distribution in \cloud\, \citep{Clark13}. In these models, the gas and dust temperatures are derived for \cloud\, as the ISRF and CRIR vary. This modelling suggest that a very high ISRF, combined with a high CRIR, reproduces well the observed discrepant dust and gas temperatures: both the ISRF and CRIR are predicted to be 1000 times higher than the values measured in the solar neighbourhood. Indeed, the models of \cite{Clark13} find that the gas and dust are not coupled, have different temperatures throughout the clump, and that \cloud\, is externally heated with a relatively cooler interior that is highly sub-structured. Despite its high mass, there appears to be very little active high-mass star formation within \cloud. The detection of a weak water maser toward it \citep{Lis94,Breen11} supports the idea that star formation may be `turning on' within this clump \citep{Lis98}. JVLA radio continuum observations reveal three compact \hii\, regions located toward the periphery of the clump; their fluxes suggest that if they lie at the distance of \cloud, then they are powered by B0.5 ZAMS stars \citep{Rodriguez13}. These \hii\, regions, however, have not yet been definitively shown to be associated with \cloud. Recent CARMA and SMA observations of N$_{2}$H$^{+}$ line and dust continuum emission show little evidence for dense material on small scales, leading \cite{Kauffmann13} to speculate that \cloud\, may lack the potential to form a cluster. A measured gas temperature of $\sim$80\,K, significantly higher than its derived dust temperature ($\sim$20\,K), led \cite{Lis01} to speculate that shocks associated with clump-clump collisions might be the dominant heating source for the gas, rather than reprocessed UV radiation. Indeed, recent H$_{2}$CO mapping of the dense gas across the CMZ measure gas kinetic temperatures ranging from 50--100\,K \citep{Ao13}, leading to the conclusion that the gas may be heated by the dissipation of turbulent energy and/or cosmic rays rather than by photon heating. This idea is also supported by other recent work that suggests that the high cosmic ray ionisation rate in the GC is responsible for the high gas temperatures (e.g. \citealp{Yusef-zadeh13}). Limited observations of SiO emission toward \cloud\, show an overall correlation between the SiO emission and high column density gas \citep{Lis01}, however, fully-sampled maps are needed to understand this apparent correspondence and to determine the details of the gas kinematics. To provide a more detailed picture of the dense and shocked gas within \cloud\, we utilise molecular line data obtained recently as part of the Millimetre Astronomy Legacy Team 90\,GHz (MALT90) survey. This survey covers many important spectral lines within the 3\,mm (90\,GHz) regime. We complement these data with additional, higher frequency APEX observations (at 230 and 345 GHz). When combined, these data can be used to measure the global physical properties and kinematics of the gas within \cloud\, and to determine whether it may be in the early stages of forming a high-mass cluster triggered by a recent close passage to Sgr\,A$^{*}$, as suggested recently by \cite{Longmore13b}.	The presence of cold dust, `hot core' chemistry, and complicated kinematics within \cloud\, is intriguing and may provide clues as to how this clump formed and whether it is on the verge of collapse. While the presence of cold dust, complex chemistry, and broad line-widths is well documented in the CMZ (e.g., \citealp{Wilson82,Martin-Pintado01}), \cloud\, is extreme compared to more typical clumps in the Galactic disk as it has a high-density, shows little star formation, and has sufficient mass to form a YMC through direct collapse. In this section we discuss the formation of \cloud\, and posit two scenarios that may explain the presence of the hot gas within \cloud\, with very different predictions for its distribution and kinematics. \subsection{The formation of \cloud\, via a close passage to Sgr\,A$^{}$} \cite{Longmore13b} suggest that the formation of \cloud\, has been triggered by the pericentre passage of a gas stream close to the bottom of the Galactic gravitational potential near Sgr\,A$^{}$, during which the gas is stretched in the orbital direction, but compressed in the direction perpendicular to its orbit. It is then argued that this compression leads to an accelerated dissipation of turbulent energy and hence triggers star formation. In this picture, \cloud\, recently passed pericentre $\sim$ 0.6\,Myrs ago and therefore should be on the verge of initiating star formation, whereas clouds like Sgr~B2 that passed pericentre 1--2 Myr ago should be exhibiting prevalent star formation. This scenario accounts for many of the observed properties of Galactic Center clouds. If this scenario accurately describes \cloud, then its molecular line emission ought to show: (1) that the dynamical time scales of its motions are comparable to the time since pericentre passage ($\sim$0.6 Myr in the model of \citealp{Molinari11}), (2) that it is elongated in the direction of its orbital motion, and (3) that there is evidence for bulk radial motions as the cloud is stretched and compressed. Each of these predictions are consistent with the observed properties of \cloud. Firstly, the implied dynamical time scale for radial motions is $\sim 0.6$ Myr, in good agreement with the estimated time since closest passage to the Galactic Center in the \cite{Molinari11} model (the optically thick lines peak at velocities red-ward of the optically thin lines by $\sim 5$ \kms\, and its effective radius is 2.9 pc, leading to a dynamical time scale, $t = r/v$, of 0.6 Myr). Second, the morphology of \cloud\, is indeed elongated, in the predicted direction along its orbit. Third, the observations clearly show bulk radial motions, as evidenced by the systematic redshift of the optically thick emission with respect to the optically thin and hot gas tracers. In the scenario where \cloud\, results from a pericentre passage, preliminary results from numerical simulations of this process (Kruijssen et al. in preparation) show that its center (within the tidal radius) will dissipate its turbulent energy at an accelerated rate and will therefore collapse, while its diffuse outer envelope (outside the tidal radius) will be stripped. Because both the collapsing center and the stripped envelope are characterised by radial motions, it is not clear whether the observed optically thick molecular lines are probing the collapsing center or the expanding envelope. Nevertheless, in either case, radial motions are clearly seen in the molecular line data presented here. The exact comparison to the numerical simulations will be presented in a future paper. \subsection{Two clumps colliding} One scenario for the presence of hot gas within \cloud\, is that it has formed recently as a result of clump-clump collisions. In this scenario, the shocks from the collisions are responsible for heating the gas. Indeed, widespread shocks, traced via SiO, are clearly associated with \cloud\, \citep{Lis01,Kauffmann13}. From large-scale CO mapping toward \cloud, \cite{Lis98} found evidence for a velocity gradient and complex kinematics within the clump that they interpret to be the superposition of two spatially overlapping components that may be interacting. In this collision/interaction scenario, the dust is not thermally coupled to the gas and, thus, has not yet had time to be heated. If true, then we would expect to see two velocity components in the dense gas tracers and a central zone of hot gas at the site of the collision. The hot/shocked gas should coincide with the position and velocity of the collision site. While our observations show two velocity components in the dense gas tracers toward the centre of the clump where presumably the components are interacting, the hot/shocked gas tracers are not isolated to this central region ( in position or velocity ). Instead, the hot/shocked gas has similar distribution and kinematics to the optically thin gas tracers and, thus, arises from similar regions across the clump. The fact that the emission from the hot/shocked gas is spread over the whole clump argues against this simple clump-clump collision scenario as a mechanism for producing the hot/shocked gas. There is no evidence for a distinct interaction zone. \subsection{A centrally condensed clump, with depletion in its cold interior} An alternative scenario is one in which the clump is a single, coherent structure, with the hot gas distributed throughout. When comparing the emission from different transitions of the same molecule, the integrated intensity images and position-velocity diagrams show that the (1--0) emission is more extended compared to the emission from the (3--2) and (4--3) transitions and isotopologues (for \hcopnt\, see Figs.~\ref{hcopoz}, \ref{hcoptt}, \ref{hcopft} and \ref{htcop}). Moreover, their spectra (Fig.~\ref{spectra-hcop}) show that the emission from the different transitions also peak at different velocities. Our data are consistent with a gradient in velocity that follows the critical density of the transitions: emission from the higher J transitions peak toward a more central velocity. This velocity shift combined with a decrease in the size of the emitting region in the higher density gas indicates a density gradient of material that is centrally condensed. Moreover, the observed anti-correlation between the dust column density and the molecular integrated intensity toward the clump's centre shows that the molecules are absent in the highest density region at the clump's center. One possible explanation for this disparity is molecular depletion in its cold interior. Thus, the absence of emission from the optically thin species toward the clump's center and systemic velocity may simply reflect depletion in the cloud's cold, dense interior. If true, then the two velocity components observed in the both the optically thin and hot/shocked gas are not tracing physically distinct clumps, but are instead simply tracing the velocity fields of a centrally condensed clump. In this scenario, the apparent presence of two distinct velocity components arises from the lack of emission in the center of a large, extended clump with a smooth velocity gradient. \subsection{Radial motions} The fact that the observed velocity field of the heated/shocked gas matches well the optically thin isotopologues implies that the emission from all of these molecules are tracing the same material. In contrast, we find that the optically thick gas is always red-shifted with respect to the optically thin/hot gas tracers (e.g. Fig.~\ref{spectra-all}). Because the optically thick lines probe the $\tau = 1$ surface, their persistent redshift with respect to the optically thin and hot gas tracers demonstrates that the cloud has radial motions and that the gas properties vary along the line of sight. Such asymmetries arise from the fact that the radial motions separate these distinct gas components since they have different radial velocities. Such an effect cannot arise from a foreground cloud; because of the large velocity gradient, such a foreground cloud would have to have exactly the same velocity gradient to absorb at precisely the right velocity to maintain the constant red-ward asymmetry. \subsubsection{P Cygni profile: an expanding, centrally condensed clump} One standard interpretation of this red-ward asymmetry is a `P Cygni' type line profile due to expanding motions (see Fig.~\ref{schematic} for a schematic). In this interpretation, the blue-shifted material arises from the outer surface of the cloud, and the red-shifted material from the cloud's interior. The brighter emission at red-shifted velocities then arises from the fact that the interior of the cloud has a significantly higher excitation temperature than the exterior portion (T$_{ex}$ inner $>$ T$_{ex}$ outer). The lower excitation temperature in the exterior layers could be due to colder temperatures in the exterior if the gas is thermalized (n $>>$ n$_{crit}$), or to sub-thermal excitation (n $<<$ n$_{crit}$). The lower excitation exterior will not only be fainter than the higher excitation temperature interior, it can also absorb line emission from the interior. Combined, these effects lead to substantial red-blue asymmetries in the line profiles of optically thick gas. \subsubsection{Baked Alaska: a collapsing, centrally condensed clump} An alternative to this `P Cygni' interpretation is a `Baked Alaska' collapse model where the clump is externally heated (T$_{K}$ outer $>$ T$_{K}$ inner). For a collapsing cloud, redshifted optically thick emission arises from the cloud's exterior and blue-shifted emission from the cloud's interior. If \cloud\, is indeed collapsing, the brighter redshifted emission arises from the fact that the exterior layers have a higher excitation temperature than the cloud's interior (see Fig.~\ref{schematic} for a schematic). In this scenario, the gas is thermalized throughout the cloud (n$>>$n$_{crit}$). This hot-exterior, cold-interior `Baked Alaska' model is consistent with the recent observations and SPH modelling of the dust and gas temperature distribution in \cloud: both indicate that the clump is externally heated \citep{Lis01,Longmore12,Clark13}. Thus, there is some evidence to support the idea that the radial motions indicate collapse, a scenario which also accounts nicely for the extremely high mass and density of the cloud. \subsection{The cluster formation potential of \cloud\, and the implications for the formation of high-mass, bound clusters} The fact that \cloud\, may be unique in the Galaxy in terms of its high density, high mass, and lack of prevalent star formation may be important for understanding how high-mass, bound clusters are formed. While the simplest idea for the formation of a high-mass cluster is through direct collapse of a large, high-mass, dense and cold molecular clump progenitor, there may in fact be multiple channels for the formation of high-mass star clusters. Rather than the simple collapse of a single molecular clump, high-mass clusters may form more slowly via gradual star formation as gas is continually accreted onto a central potential well from a more extend reservoir of material. While plausible, the observed lack of age spreads in young high-mass clusters argues against this formation scenario \citep{Clark05,Negueruela10,Kudryavtseva12}. Massive clusters may also form through mergers of smaller groups of stars that have already formed but are part of a common potential well (e.g., \citealp{McMillan07,Allison09}). The observed gas morphology and kinematics suggest that \cloud\, is indeed a single, coherent high-mass dense, clump that may be highly fragmented on small spatial scales. As such, it has the potential to form a high-mass cluster. The puzzle, however, is how a dense, 10$^{5}$ \Msun\, clump forms without rapidly producing stars. For \cloud, its location in the CMZ may provide the clue: the increased turbulence in the CMZ may in fact inhibit star formation below a column density threshold which is much higher in the CMZ than in the Galactic plane \citep{Kruijssen13a}. However, the fact that some of the most massive Galactic YMCs are located outside of the Galactic centre region (e.g., Westerlund 1, Glimpse-C01, NGC 3603, RSGCs) suggests that the unique conditions within the CMZ are not essential for their formation. High-mass clusters are thought to form quickly, perhaps in less than a few Myrs. Recent observations of NGC 3603 and Westerlund 1 suggests that the age spread within the stellar population may be $<$0.4 Myr \citep{Clark05,Negueruela10,Kudryavtseva12}. The fact that \cloud\, is one of a handful of clumps with sufficient mass to form a YMC may support this fast formation scenario; at any given time the number of high-mass cluster precursors within the Galaxy will be limited to just a few \citep{Longmore12}. If true, then perhaps we are catching \cloud\, at a very special time, immediately before the formation of the high-mass cluster. In this case, we would expect the gas and dust to be highly fragmented, the fragments being the precursors to the individual stars or sub-clusters. Recent models do predict that \cloud\, should form a bound cluster through hierarchical fragmentation \citep{Kruijssen12b}. While our data show evidence for fragmentation, higher-angular resolution data of both the optically thin and shocked gas tracers down to small scales ($<$ 0.1pc) are clearly needed to definitively determine whether or not \cloud\, with give rise to a cluster in the future.	14	3	1403.0996	Despite their importance as stellar nurseries and the building blocks of galaxies, very little is known about the formation of the highest mass clusters. The dense clump G0.253+0.016 represents an example of a clump that may form an Arches-like, high-mass cluster. Here we present molecular line maps toward G0.253+0.016 taken as part of the MALT90 molecular line survey, complemented with APEX observations. Combined, these data reveal the global physical properties and kinematics of G0.253+0.016. Recent Herschel data show that while the dust temperature is low (~19 K) toward its center, the dust temperature on the exterior is higher (~27 K) due to external heating. Our new molecular line data reveal that, overall, the morphology of dense gas detected toward G0.253+0.016 matches its IR extinction and dust continuum emission very well. An anticorrelation between the dust and gas column densities toward its center indicates that the clump is centrally condensed with a cold, dense interior in which the molecular gas is chemically depleted. The velocity field shows a strong gradient along the clump's major axis, with the blueshifted side at a higher Galactic longitude. The optically thick gas tracers are systematically redshifted with respect to the optically thin and hot gas tracers, indicating radial motions. The gas kinematics and line ratios support the recently proposed scenario in which G0.253+0.016 results from a tidal compression during a recent pericenter passage near Sgr A*. Because G0.253+0.016 represents an excellent example of a clump that may form a high-mass cluster, its detailed study should reveal a wealth of knowledge about the early stages of cluster formation.	false	[ "cluster formation", "dense gas", "the highest mass clusters", "molecular line maps", "external heating", "line ratios", "APEX observations", "continuum emission", "G0.253", "dust", "*", "Sgr A", "knowledge", "a high-mass cluster", "radial motions", "Recent Herschel data", "a higher Galactic longitude", "K", "+", "the MALT90 molecular line survey" ]	11.2767	10.2561	-1
481728	[ "Parsons, R. D.", "Hinton, J. A." ]	2014APh....56...26P	[ "A Monte Carlo template based analysis for air-Cherenkov arrays" ]	137	[ "Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany", "Department of Physics and Astronomy, The University of Leicester, University Road, Leicester LE1 7RH, United Kingdom" ]	[ "2015ICRC...34..777L", "2015ICRC...34..801D", "2015ICRC...34..873R", "2015JPhG...42i5201S", "2015arXiv150900794M", "2015arXiv150901429M", "2015arXiv150905791M", "2015arXiv151005675H", "2016MNRAS.458.2813V", "2017A&A...597A.115H", "2017A&A...600A..89H", "2017A&A...606A..59H", "2017AIPC.1792d0005P", "2017AIPC.1792d0017B", "2017AIPC.1792d0033M", "2017AIPC.1792d0035M", "2017AIPC.1792d0040S", "2017AIPC.1792e0023Z", "2017AIPC.1792e0029C", "2017AIPC.1792f0013S", "2017ApJ...850L..22A", "2017IAUS..331..325S", "2017ICRC...35..540O", "2017ICRC...35..627C", "2017ICRC...35..645G", "2017ICRC...35..654Z", "2017ICRC...35..655Z", "2017ICRC...35..675R", "2017ICRC...35..676H", "2017ICRC...35..707M", "2017ICRC...35..711S", "2017ICRC...35..789C", "2017ICRC...35..811C", "2017ICRC...35..837S", "2017ICRC...35..846M", "2017ICRC...35..905O", "2017NPPP..291...30B", "2017PhDT.......254A", "2017arXiv171106298L", "2018A&A...612A...5H", "2018A&A...612A...6H", "2018A&A...612A..12H", "2018A&A...619A..71H", "2018APh....97....1S", "2018JCAP...11..037A", "2018MNRAS.476.4187H", "2018NIMPA.907...31H", "2018PhRvD..98b2009A", "2018PhRvL.120t1101A", "2018Sci...361.1378I", "2018arXiv181004516H", "2019A&A...621A.116H", "2019A&A...626A..57H", "2019A&A...627A.159H", "2019APh...105...44S", "2019APh...111...35A", "2019ASPC..523...75A", "2019ApJ...881..134A", "2019EPJC...79..427S", "2019Galax...7...21W", "2019ICRC...36..518G", "2019ICRC...36..732M", "2019ICRC...36..753N", "2019ICRC...36..787S", "2019JCAP...01..012J", "2019JCAP...06..042L", "2019MNRAS.482.3011A", "2019MNRAS.486.3886H", "2019Natur.575..464A", "2019hepr.confE..27W", "2020A&A...633A.102H", "2020A&A...633A.162H", "2020A&A...639A..42A", "2020APh...114...92S", "2020APh...12302491H", "2020ApJ...894L..16A", "2020EPJC...80..363P", "2020MNRAS.494.5590A", "2020NatAs...4..167H", "2020Natur.582..356H", "2020PhRvD.102f2001A", "2021A&A...648A..23H", "2021A&A...655A...7H", "2021APh...12402508D", "2021APh...12902579S", "2021ApJ...911L..11E", "2021ApJ...917....6A", "2021ApJ...918...17A", "2021ApJ...919..106A", "2021ApJ...923..109A", "2021ApJ...923..241A", "2021EPJC...81.1101O", "2021JInst..16P7050S", "2021PhRvD.103b3011R", "2021PhRvD.103j2002A", "2021Sci...372.1081H", "2021arXiv210801331G", "2021arXiv210802282A", "2022A&A...662A..65H", "2022A&A...666A.124A", "2022EPJC...82.1118O", "2022Galax..10...21S", "2022MNRAS.515.1365C", "2022MNRAS.517.4736A", "2022PhRvL.129k1101A", "2022Sci...376...77H", "2022Univ....8...90D", "2022arXiv220904045D", "2022icrc.confE.780M", "2022icrc.confE.805P", "2022icrc.confE.889C", "2022icrc.confE.955D", "2023A&A...672A.103H", "2023A&A...673A.148H", "2023A&A...675A.138H", "2023A&A...677A.141L", "2023APh...14802817G", "2023ApJ...946L..27A", "2023ApJ...950L..16A", "2023ApJ...952L..38A", "2023ApJ...954...70A", "2023JCAP...04..040A", "2023NatAs...7.1341H", "2023arXiv230602460C", "2023arXiv230808969R", "2023arXiv230812732D", "2023hxga.book..137M", "2024A&A...683A..70H", "2024A&A...686A.149H", "2024A&C....4600793D", "2024APh...15802937I", "2024RDTM..tmp....7L", "2024Sci...383..402H", "2024Univ...10..100L", "2024arXiv240312608A", "2024arXiv240417623T", "2024arXiv240617502S" ]	[ "astronomy" ]	32	[ "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1957SoEn....1...16D", "1975CoPhC..10..343J", "1983ApJ...272..317L", "1998NIMPA.416..278B", "1998NIMPA.416..425L", "1998cmcc.book.....H", "1999APh....12..135H", "2004NewAR..48..331H", "2006APh....25..195L", "2008APh....30..149B", "2009APh....31..383O", "2009APh....32..231D", "2011APh....34..858B", "2013APh....43..171B" ]	[ "10.1016/j.astropartphys.2014.03.002", "10.48550/arXiv.1403.2993" ]	1403	1403.2993_arXiv.txt	Ground-based gamma-ray astronomy exploits the air shower produced by the interaction of a primary gamma ray in the Earth's atmosphere. For gamma-rays of approximately $10^{11} - 10^{14}$ eV (VHE) few air-shower particles reach ground level. However, sufficiently high-energy secondary particles will emit Cherenkov radiation, resulting in illumination of a $\sim$10$^{5}$ m$^{2}$ patch of ground for a few nanoseconds. The Imaging Atmospheric Cherenkov Technique (IACT) is based on the observation of this emission by a number of large reflecting telescopes, placed within the light pool with a typical spacing of $\sim$100~m. The Cherenkov light is focussed onto ultra-fast, but typically rather coarsely pixelated, cameras, resulting in roughly elliptical images of the shower emission in multiple telescopes. H.E.S.S. \citep{Hinton2004} is an array of four such 100\,m$^{2}$ reflecting telescopes, at 1800\,m altitude in the Khomas Highlands of Namibia, each with a 960 pixel camera viewing a 5$^\circ$ diameter patch of sky. The process of reconstruction/estimation of the properties of the primary photon (direction and energy), and the rejection of events likely to belong to the background of hadronic showers, makes use of image information from each telescope. Traditionally this event reconstruction is performed using the \textit{Hillas} parameters of the camera images~\cite{Hillas1985}, derived after an image cleaning step (described in \cite{Aharonian2006}). The Hillas parameters are the moments of the camera image which, given the approximately elliptical nature of typical camera images, already capture much of the available image information. In the most commonly used stereoscopic reconstruction method the major axes of images are calculated in a common camera reference frame and the intersection points of all axes found. A weighted average (based on image amplitude and the angle between the axes) is then taken of all crossing points to provide an estimate of the arrival direction of the primary gamma-ray. A similar procedure involving the intersections of the directions between the image centroid and the optical axis is then performed in a common plane perpendicular to the pointing direction, to determine the shower impact point on the ground. Although relatively good angular resolution ($\sim$0.1$^\circ$ with H.E.S.S.) can be reached with such a reconstruction procedure~\citep{Aharonian2006}, significant additional information can be extracted from the recorded images in a typical event, resulting in improved performance. Additionally a simultaneous fit procedure between all telescopes can help to ensure a consistent result is found between all telescopes, rather than the independent analysis used in Hillas reconstruction. An image template fitting procedure was pioneered for the CAT telescope \citep{Barrau1998} \cite{LeBohec1998} and re-implemented and improved upon for H.E.S.S. \cite{deNaurois2009}. These methods begin with the creation of a semi-analytical model of air-shower development and IACT response, and use this model to generate the expected shower image for a given set of shower parameters (primary energy and direction, and also impact distance of the shower from the telescope). The template library can then be compared to the images recorded in a given event, and, by means of a multi-dimensional fit procedure, the best-fit shower parameters determined. An alternative method of shower fitting has also been developed for use with H.E.S.S. data, fitting the pixel intensities of the camera image with the expectation from a simpler analytical 3 dimensional gaussian air shower model (3D model) \citep{Lemoine2006,Becherini2011}. One of the major problems with the \textit{model} and 3D model analyses is the difficulty of describing the air shower behaviour at high energies ($>$10\,TeV). Above 10\,TeV a large number of particles reach ground level, resulting in large fluctuations which are difficult to reproduce with the aforementioned approaches. This difficultly in reproducing energetic air showers results in poor event reconstruction above 10\,TeV, typically leading to a rapid drop in effective collection area. With these approaches it is also difficult to account for instrumental effects, such as the optical point spread function of the telescopes or the limited readout window of the camera, so approximations of these effects must be made. Inevitably, making these kind of assumptions will limit the quality of the model fit, reducing the accuracy of the analysis. One way to solve these problems with high energy events is to instead produce templates by the use of detailed Monte Carlo (MC) based air shower simulations. Such an approach requires no analytical approximations to be made within the air shower simulation, and therefore takes better into account the large fluctuations arising from particles reaching ground level. Additionally, a more accurate optics and electronics simulation can be performed on the simulated Cherenkov photons, resulting a more realistic representation of the telescope behaviour. The MC approach is a robust and general image template generation method, not requiring any adaptation to the model for different telescope types, instead existing Monte Carlo telescope configurations can be adapted to produce templates. Below we present an attempt to improve the accuracy of the template generation, by the use of a more accurate Monte Carlo simulation based air shower model, combined with a ray-tracing telescope simulation and demonstrate the resultant improvement in the performance of the analysis. \begin{figure}[t] \begin{center} \includegraphics[width=0.99\textwidth]{templates.eps} \caption{Image template histograms for a 1\,TeV primary gamma-ray at a core distance of 20\,m (top left), 100\,m (top right) and 200\,m (bottom left) at the expected shower maximum (300 g\,cm$^{-2}$). The bottom right plot shows the shower template at a core distance of 100\,m and a X$_{\rm max}$ of 400 g\,cm$^{-2}$. The $z$-axis is shown in units of photoelectrons per square degree, and the $x-$ and $y-$axes are in degrees.} \label{fig-templates} \end{center} \end{figure}	We have demonstrated a high performance likelihood-based reconstruction method for the H.E.S.S. telescope, based on expected image templates generated from Monte Carlo air shower simulations. This technique has been proven to be extremely successful, reconstructing events with significantly more accuracy than the standard reconstruction algorithm. Some improvement has also been demonstrated over the \textit{model} reconstruction, especially at the highest energies where semi-analytical modelling of the air shower becomes difficult. Use of the \impact{} reconstruction is able to provide an sensitivity improvement of around a factor of 2 in observation time over traditional \textit{Hillas-based} reconstruction for point source observations. Additionally the improvement in angular resolution offers more detailed imaging of extended sources. Although some performance improvements have been shown over the existing, semi-analytical model based approach to template generation, the major advantage of production via MC simulation is the robustness of this approach against changes in telescope hardware. For example the H.E.S.S. observatory has recently commissioned a fifth 26\,m diameter telescope (CT5) at the centre of the array, which requires the production of a new set of image templates. In order to generate these new templates one can simply re-run the template generation step using the CT5 simulation configuration, whereas using a model based approach the model must be changed to accurately reflect the new hardware. This robustness will be especially important for the next generation of Cherenkov telescopes, such as the Cherenkov Telescope Array, which plan to use multiple telescope types within the same array. Additionally, this template generation mechanism can in principle be implemented to allow reconstruction of any particle type. In the case of protons the large shower to shower fluctuations present may make this difficult, however for other particles such as electrons or Iron nuclei it may be possible to reconstruct and separate events using this method.	14	3	1403.2993	We present a high-performance event reconstruction algorithm: an Image Pixel-wise fit for Atmospheric Cherenkov Telescopes (ImPACT). The reconstruction algorithm is based around the likelihood fitting of camera pixel amplitudes to an expected image template. A maximum likelihood fit is performed to find the best-fit shower parameters. A related reconstruction algorithm has already been shown to provide significant improvements over traditional reconstruction for both the CAT and H.E.S.S. experiments. We demonstrate a significant improvement to the template generation step of the procedure, by the use of a full Monte Carlo air shower simulation in combination with a ray-tracing optics simulation to more accurately model the expected camera images. This reconstruction step is combined with an MVA-based background rejection.	false	[ "camera pixel amplitudes", "Monte Carlo", "a full Monte Carlo air shower simulation", "an expected image template", "the template generation step", "Atmospheric Cherenkov Telescopes", "the expected camera images", "H.E.S.S. experiments", "combination", "traditional reconstruction", "significant improvements", "a ray-tracing optics simulation", "This reconstruction step", "MVA", "H.E.S.S.", "an MVA-based background rejection", "CAT", "A maximum likelihood fit", "A related reconstruction algorithm", "the procedure" ]	7.189602	0.540763	13
483012	[ "Miller, J. M.", "Bachetti, M.", "Barret, D.", "Harrison, F. A.", "Fabian, A. C.", "Webb, N. A.", "Walton, D. J.", "Rana, V." ]	2014ApJ...785L...7M	[ "Patchy Accretion Disks in Ultra-luminous X-Ray Sources" ]	20	[ "Department of Astronomy, University of Michigan, 500 Church Street, Ann Arbor, MI 48109-1042, USA", "Universite de Toulouse, UPS-OMP, IRAP, F- 31100 Toulouse, France; CNRS, Intitut de Recherche in Astrophysique et Planetologie, 9 Av. Colonel Roche, BP 44346, F-31028 Toulouse, Cedex 4, France", "Universite de Toulouse, UPS-OMP, IRAP, F- 31100 Toulouse, France; CNRS, Intitut de Recherche in Astrophysique et Planetologie, 9 Av. Colonel Roche, BP 44346, F-31028 Toulouse, Cedex 4, France", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Universite de Toulouse, UPS-OMP, IRAP, F- 31100 Toulouse, France; CNRS, Intitut de Recherche in Astrophysique et Planetologie, 9 Av. Colonel Roche, BP 44346, F-31028 Toulouse, Cedex 4, France", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA" ]	[ "2014ApJ...793...21W", "2015ApJ...799..122W", "2015ApJ...806...65W", "2015ApJ...808...64M", "2015JApA...36..233Z", "2015MNRAS.447.3243M", "2015MNRAS.448.1153P", "2016AN....337..534R", "2016ApJ...826L..26W", "2016MNRAS.456.1859U", "2016MNRAS.460.4417L", "2017ApJ...839...46S", "2017ApJ...839..105W", "2018ApJ...856..128W", "2018MNRAS.473.4360W", "2018MNRAS.474.1214Y", "2018MNRAS.476.2045Y", "2018MNRAS.479.4271P", "2019MNRAS.488..451M", "2020MNRAS.494.6012W" ]	[ "astronomy" ]	7	[ "accretion", "accretion disks", "black hole physics", "X-rays: binaries", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1979ApJ...228..939C", "1984PASJ...36..741M", "1988ngc..book.....T", "1994ApJ...434..570T", "1996ASPC..101...17A", "1998MNRAS.297..929G", "1999ApJ...512..892T", "2000ApJ...542..914W", "2000MNRAS.313..193M", "2001ApJ...551..897B", "2001ApJ...552L.109K", "2001MNRAS.321..549M", "2002ApJ...568L..97B", "2003ApJ...585L..37M", "2003ApJ...586L..61S", "2003ApJ...587..356I", "2003MNRAS.341.1041A", "2004ApJ...614L.117M", "2006MNRAS.368..397S", "2007ApJ...660..580S", "2007ApJ...670...92G", "2009MNRAS.397.1836G", "2009MNRAS.398.1450K", "2009Natur.459..540F", "2009PASP..121.1279S", "2011AN....332..330S", "2011MNRAS.413.1623D", "2011MNRAS.417..464M", "2012ApJ...753..139P", "2012MNRAS.420.1107P", "2012MNRAS.422..990K", "2012MNRAS.423.1154S", "2012MNRAS.426..473W", "2012MNRAS.426L..71D", "2013ApJ...764...93P", "2013ApJ...770..103H", "2013ApJ...773L...9W", "2013ApJ...775...75M", "2013ApJ...776L..36M", "2013ApJ...778..163B", "2013ApJ...779..148W", "2013MNRAS.430.1408K", "2013MNRAS.433.1023G", "2014MNRAS.438L..51M", "2014MNRAS.439L...1C" ]	[ "10.1088/2041-8205/785/1/L7", "10.48550/arXiv.1403.1769" ]	1403	1403.1769_arXiv.txt	Ultra-luminous X-ray (ULXs) are accretion-powered sources that appear to violate the isotropic Eddington limit for a fiducial $M = 10~M_{\odot}$ black hole. In practice, a better lower limit is $L_{X} \geq 2\times 10^{39}~{\rm erg}~ {\rm s}^{-1}$ (Irwin, Athey, \& Bregman 2003). Most or all of the sources that barely qualify as ULXs are likely to be stellar-mass black holes similar to those known in the Milky Way. The set of ULXs with $L_{x} \geq 10^{40}~ {\rm erg}~ {\rm s}^{-1}$ is far more interesting, as these sources more strongly indicate super-Eddington accretion or elevated black hole masses. Only a small number of these extreme ULXs are found at distances that permit sensitive spectra in reasonable observation times (note that Sutton et al.\ 2012 and Gladstone 2013 define ``extreme'' differently). Yet, detailed studies have revealed multiple components in the spectra of these ULXs, including soft disk--like components. The low temperature values typical of these components ($kT \simeq 0.2$~keV) may indicate accretion onto more massive black holes (e.g. $10^{2-3}~ M_{\odot}$; Miller et al.\ 2003, 2004, 2013). Mass estimates depend strongly on numerous assumptions (e.g. Soria 2011, Miller et al.\ 2013), not least the idea that black holes in ULXs accrete in a mode that is observed in sub-Eddington stellar-mass black holes and in AGN. However, ULXs may show spectral phenomena that stellar-mass black holes and AGN may not. In particular, deep {\it XMM-Newton} spectra show evidence of a roll-over in the 5--10~keV band. This peculiar hard component can also be modeled as Comptonization of low temperature photons (e.g. Stobbart et al.\ 2006, Gladstone et al.\ 2009). In constrast to the hot, optically--thin Comptonization regions inferred in other black holes, the roll-over in ULXs requires a cool, optically--thick region ($kT_e \simeq 2$~keV; $\tau \simeq 10$). Unless the corona is heated locally and/or powered magnetically -- in some manner that would prevent substantial heating -- it is easy to show that such a cool corona must be huge in order to generate the bulk of the observed $L_{X} \geq 10^{40}~ {\rm erg}~ {\rm s}^{-1}$ that is observed (see, e.g., Merloni \& Fabian 2001 concerning thermal and magnetic coronae). Not only is it difficult to envisage a very large region that can be characterized by a single temperature, but the outer part of the putative corona may only be marginally bound to the black hole. Some extreme ULXs have been detected up to 40~keV in recent {\it NuSTAR} (Harrison et al.\ 2013) observations, and the roll--over found in prior 0.3--10.0~keV spectra is strongly confirmed (e.g Walton et al.\ 2013a, 2014; Bachetti et al.\ 2013; Rana et al.\ 2014). Low--temperature Comptonization again provides acceptable fits, but again is not a unique description, and the problems of such coronae remain. Putative Comptonization components with such low electron temperatures and high optical depths, are fairly similar to a blackbody (e.g. Stobbart et al.\ 2006; Kajava et al.\ 2012; Middleton et al.\ 2011). Miller et al.\ (2013) recently reported positive correlations between the temperature and luminosity of the cool ($kT \simeq 0.2$~keV) disk--like components in ULXs, suggesting that a disk interpretation may be also be correct or that component (less stringent data selection and modeling previously led to different conclusions; see Kajava \& Poutanen 2009; Pintore \& Zampieri 2012). In view of that result -- and the blackbody--like nature of the harder component that rolls-over at high energy -- Miller et al.\ (2013) suggested that both components may originate in the disk, and that ULX spectra may be consistent with ``patchy'' or inhomogeneous disks (see Dexter \& Quataert 2012; also see Begelman 2002, Dotan \& Shaviv 2011). This Letter explicitly examines the possibility of inhomogeneous disks in extreme ULXs.	Following the suggestion of patchy or inhomogeneous accretion disks in ULXs in Miller et al.\ (2013), we have specifically examined the ability of phenomenological descriptions of a ``patchy'' accretion disk (e.g. Dexter \& Quataert 2012) to describe the spectra of extreme ULXs. At least in the case of NGC 1313 X-1, our fits to two epochs of joint {\it XMM-Newton} and {\it NuSTAR} spectra suggest that a patchy disk is potentially an excellent description of the observed spectra. Importantly, such models likely avoid the need for very large but cool Comptonization regions implied by recent spectral models (e.g. Gladstone et al.\ 2009). Theoretical work has found evidence of photon bubble instabilities (Gammie 1998, Begelman 2001, 2002) that might naturally give rise to a patchy temperature profile and locally super-Eddington flux levels. A patchy disk need not be a unique signature of a photon bubble instability; other instabilities or mechanisms might be able to produce a similar effect. It is possible that a two--temperature or patchy disk is the natural signature of near-Eddington or super-Eddington accretion, and that a compelling correspondence between theory and observations has been identified. But, it is important to consider caveats and means of testing this scenario. King et al.\ (2001) described a model for super-Eddington accretion in ULXs wherein a funnel-like geometry develops in the inner disk, mechanically beaming some radiation and foiling luminosity estimates based on isotropic emission. The hot, small region implied by our two-temperature models could correspond to the inner wall of a funnel. However, recent observations of the extreme ULX NGC 5408 X-1 have revealed X-ray flux dips typical of disks viewed at {\it high} inclination (e.g. Pasham \& Strohmayer 2013, Grise et al.\ 2013). Moreover, recent radio observations of Holmberg II X-1 -- which can also be fit with this spectral model -- reveal three components (Cseh et al.\ 2014), potentially indicating a jet system viewed at high inclination. It is possible that the hot disk emission corresponds to patches rather than to a line of sight that peers into a funnel. If ULXs are not all viewed at low inclination angles, it is possible that their radiation is nearly isotropic, and that ULXs with $L_{X} \geq 10^{40}~ {\rm erg}~ {\rm s}^{-1}$ harbor black holes with masses above the range known in Galactic X-ray binaries. This may be supported by recent evidence that cool thermal components in ULXs appear to be broadly consistent with the $L \propto T^{4}$ trend expected for standard thin disk accretion (Miller et al.\ 2013). However, this does not explain why the high energy spectra of extreme ULXs differ so markedly from stellar-mass black holes and most AGN (Bachetti et al.\ 2013, Walton et al.\ 2013a, 2014). A patchy disk model may be able to account for both spectral features. A wider set of sources and states must be sampled in order to better test the possibility of patchy disks in ULXs. Variability below 10~keV might be ascribed to changes in the characteristic temperatures -- and relative emitting areas -- of cooler and hotter disk regions. Such changes might be linked to the mass accretion rate through the disk. If a mechanism like the photon bubble instability is at work, spectral variations might also be tied to changes in the disk magnetic field (which could itself be linked to the mass accretion rate). Spectral variability above approximately 10~keV might be driven partially by changes in an hotter disk component, but also changes in a corona. New observations of Holmberg IX X-1 in different states may be consistent with a patchy accretion disk model (Walton et al.\ 2014). It is possible that different flux ratios from the larger disk and hot patches could account for differences between very soft sources (such as NGC 5408 X-1), and harder sources (like Holmberg IX X-1). If patchy disks power ULXs to super-Eddington luminosities, a strong wind might be driven (but see Begelman 2001). However, winds have not been conclusively detected in ULXs. Indeed, narrow absorption lines with strengths comparable to features detected in Galactic sources are ruled out at high confidence levels in deep spectra (Walton et al.\ 2012, 2013b; Pasham \& Strohmayer 2012). Emission lines from an outflow perpendicular to the line of sight with equivalent widths comparable to SS 433 ($few \times 100$~eV; see, e.g., Marshall et al.\ 2013) are also ruled out. Potential low-energy lines in NGC 5408 X-1 have been discussed as having a wind origin (Middleton et al.\ 2014), but the lines may also be described in terms of diffuse emission with a constant flux (e.g. Miller et al.\ 2013). If the disk is only super-Eddington in localized regions -- the hot patches -- then radiation pressure would only drive winds in localized regions. Moreover, even if the temperature profile is more shallow than the $T \propto R^{-3/4}$ derived for standard thin disks, the emission and wind regions will be concentrated close to the black hole. Any winds will therefore be highly ionized and -- assuming the wind must retain some of the angular momentum of its launching point -- highly smeared by rotation. Such effects would certainly inhibit the detection of a wind via {\it narrow} absorption lines. It is possible that strong QPOs could be generated via a patchy disk profile. This might also cause the QPOs to be stronger at the higher energies that correspond to the hot patches; this would match observed relations in more standard X-ray binaries. QPOs are observed in some extreme ULXs (e.g. M82 X-1, NGC 5408 X-1; Strohmayer \& Mushotzky 2003, Strohmayer et al.\ 2007). Recent simulations suggest that patchy disk profiles do not necessarily lead to QPO production (Armitage \& Reynolds 2003), but additional theoretical work and observational tests may be required. Relativistically--blurred disk reflection is the most likely explanation for the bulk of the ``soft excess'' observed in NLSy1s (Fabian et al.\ 2009, Kara et al.\ 2013). Yet, even in a case such as 1H 0707$-$495 -- wherein the disk reflection spectrum is particularly strong and clear -- a very soft blackbody component is still required at the edge of the band pass (Fabian et al.\ 2009). It is possible that this emission represents the high energy tail of a patchy accetion disk that puts out more power in its lower--temperature UV/EUV component. It is interesting to note that some sources identified by Greene \& Ho (2007) appear to have similar X-ray spectra, and ``slim'' or super-Eddington disk models may apply in some cases. \vspace{0.2in} JMM acknowledges helpful conversations with Mitch Begelman, Jason Dexter, and Richard Mushotzky. This work was supported under NASA Contract No. NNG08FD60C, and made use of data from the NuSTAR mission, a project led by the California Institute of Technology, managed by the Jet Propulsion Laboratory, and funded by NASA.	14	3	1403.1769	The X-ray spectra of the most extreme ultra-luminous X-ray sources—those with L >= 10<SUP>40</SUP> erg s<SUP>-1</SUP>—remain something of a mystery. Spectral roll-over in the 5-10 keV band was originally detected in the deepest XMM-Newton observations of the brightest sources; this is confirmed in subsequent NuSTAR spectra. This emission can be modeled via Comptonization, but with low electron temperatures (kT<SUB>e</SUB> ~= 2 keV) and high optical depths (τ ~= 10) that pose numerous difficulties. Moreover, evidence of cooler thermal emission that can be fit with thin disk models persists, even in fits to joint XMM-Newton and NuSTAR observations. Using NGC 1313 X-1 as a test case, we show that a patchy disk with a multiple temperature profile may provide an excellent description of such spectra. In principle, a number of patches within a cool disk might emit over a range of temperatures, but the data only require a two-temperature profile plus standard Comptonization, or three distinct blackbody components. A mechanism such as the photon bubble instability may naturally give rise to a patchy disk profile, and could give rise to super-Eddington luminosities. It is possible, then, that a patchy disk (rather than a disk with a standard single-temperature profile) might be a hallmark of accretion disks close to or above the Eddington limit. We discuss further tests of this picture and potential implications for sources such as narrow-line Seyfert-1 galaxies and other low-mass active galactic nuclei.	false	[ "subsequent NuSTAR spectra", "such spectra", "super-Eddington luminosities", "NuSTAR observations", "low electron temperatures", "accretion disks", "sources", "thin disk models", "numerous difficulties", "temperatures", "the most extreme ultra-luminous X-ray sources", "NuSTAR", "high optical depths", "other low-mass active galactic nuclei", "standard Comptonization", "The X-ray spectra", "a patchy disk profile", "a multiple temperature profile", "L", "gt;=" ]	6.648159	6.904356	22
437501	[ "Chen, Yan-Ping", "Trager, S. C.", "Peletier, R. F.", "Lançon, A.", "Vazdekis, A.", "Prugniel, Ph.", "Silva, D. R.", "Gonneau, A." ]	2014A&A...565A.117C	[ "The X-shooter Spectral Library (XSL). I. DR1: Near-ultraviolet through optical spectra from the first year of the survey" ]	98	[ "Kapteyn Astronomical Institute, University of Groningen, PO BOX 800, 9700 AV, Groningen, The Netherlands", "Kapteyn Astronomical Institute, University of Groningen, PO BOX 800, 9700 AV, Groningen, The Netherlands", "Kapteyn Astronomical Institute, University of Groningen, PO BOX 800, 9700 AV, Groningen, The Netherlands", "Observatoire Astronomique, 11 rue de l'Université, 67000, Strasbourg, France", "Instituto de Astrofísica de Canarias (IAC), 38200 La Laguna, Tenerife, Spain; Departamento de Astrofísica, Universidad de La Laguna, 38205, Tenerife, Spain", "Université de Lyon, Université Lyon 1, 69622, Villeurbanne, 69561, Saint-Genis Laval, France", "National Optical Astronomy Observatory, 950 North Cherry Avenue, Tucson, AZ 85719, USA", "Observatoire Astronomique, 11 rue de l'Université, 67000, Strasbourg, France" ]	[ "2014A&A...567A...6P", "2014A&A...568A...9M", "2014A&A...572A..13S", "2014A&A...572A..47R", "2015A&A...579L...6B", "2015A&A...582A..49H", "2015ASPC..497..113G", "2015MNRAS.449.1177V", "2015MNRAS.451.3237W", "2015MNRAS.451.3504R", "2015MNRAS.452.2434S", "2015MNRAS.453.3510H", "2015arXiv151200688Z", "2016A&A...585A..64S", "2016A&A...585A..87S", "2016A&A...589A..36G", "2016A&A...589A..73R", "2016ASPC..507..121H", "2016MNRAS.455.4467M", "2016MNRAS.462.1461M", "2016MNRAS.463..886D", "2016PhDT........93C", "2016SPIE.9908E..1GD", "2017A&A...597A..82N", "2017A&A...601A.141G", "2017A&A...603A.119H", "2017AJ....154...29K", "2017ASInC..14...21G", "2017ApJ...846..132H", "2017ApJ...850...34B", "2017CaJPh..95..840M", "2017MNRAS.464..194F", "2017MNRAS.472..361R", "2017NewA...56...71C", "2018A&A...615A..24W", "2018A&A...615A.115K", "2018A&A...617A..37C", "2018A&A...619A.134H", "2018ApJ...861...67C", "2018ApJ...867..110B", "2018MNRAS.473..826H", "2018MNRAS.474.2116D", "2018MNRAS.476.5189H", "2018MNRAS.480..629M", "2018MNRAS.480.4766W", "2018arXiv181102841L", "2019A&A...621A..60F", "2019A&A...627A.138A", "2019A&A...629A.100I", "2019A&A...632A.105C", "2019ApJ...877..102K", "2019ApJ...878..134K", "2019ApJ...878..147K", "2019ApJ...883..175Y", "2019IAUS..343..309L", "2019MNRAS.484.4619G", "2019MNRAS.486.3228R", "2019Msngr.178...17I", "2019RNAAS...3....3K", "2020A&A...634A.114C", "2020A&A...634A.133G", "2020A&A...640A..92M", "2020A&A...641A..44M", "2020ApJ...899...62C", "2020ApJ...901..161K", "2020MNRAS.491.2025C", "2020MNRAS.492..782W", "2020MNRAS.492.2177K", "2020MNRAS.492.4216S", "2020RAA....20..148C", "2021A&A...649A..97L", "2021AJ....161...52L", "2021MNRAS.501.4596G", "2021MNRAS.505.4496G", "2021MNRAS.506.1557W", "2022A&A...660A..34V", "2022A&A...661A..50V", "2022A&A...668A..21L", "2022AJ....163...56I", "2022AJ....164...89H", "2022ApJS..261....4O", "2022ApJS..261...20M", "2022MNRAS.509.4308H", "2022MNRAS.513.5920F", "2023A&A...672A.166S", "2023A&A...672A.181S", "2023A&A...674A...8F", "2023AJ....166..236D", "2023ApJ...956...60C", "2023ApJS..266...11B", "2023MNRAS.519.2281L", "2023MNRAS.521.2745S", "2023MNRAS.526.3273C", "2024A&A...682A.133K", "2024ApJ...965..106Y", "2024MNRAS.530.3651B", "2024MNRAS.tmp.1481C", "2024arXiv240201631T" ]	[ "astronomy" ]	21	[ "stars: abundances", "stars: fundamental parameters", "stars: AGB and post-AGB", "stars: atmospheres", "galaxies: stellar content", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]	[ "1966CoLPL...4...99J", "1971ROAn....7.....C", "1973AJ.....78..959L", "1973MmRAS..77..223C", "1983AJ.....88..439L", "1983bscs.book.....H", "1984AJ.....89..636B", "1986PASP...98..609H", "1987A&A...186....1G", "1988ARA&A..26...51F", "1989PhDT.......149P", "1990AJ.....99..784H", "1990PASP..102.1181B", "1991ApJ...367..126C", "1991bsc..book.....H", "1992A&AS...96..593L", "1992ApJ...398...69W", "1993MNRAS.262..650D", "1993PhDT.......172G", "1993sssp.book.....K", "1994AJ....107..513B", "1994ApJS...94..687W", "1995PASP..107..945F", "1996A&AS..117..227A", "1996AJ....111.1748F", "1996PASP..108..996L", "1997A&A...326..950F", "1997ApJS..111..377W", "1998ApJS..119..121L", "1998MNRAS.300..872M", "1998PASP..110..863P", "1999A&AS..140..261A", "1999ApJS..123....3L", "2000A&AS..146..217L", "2000AJ....119.1645T", "2001A&A...369.1048P", "2001MNRAS.326..959C", "2001PASP..113.1420V", "2002A&A...393..167L", "2003A&A...402..433L", "2003A&A...406...51C", "2003MNRAS.344.1000B", "2003Msngr.114...10B", "2004A&A...425..881L", "2004ApJS..152..251V", "2004PASP..116..138C", "2005A&A...442.1127M", "2005A&A...443..143G", "2005A&A...443..735C", "2005ApJ...621..673T", "2005ApJ...628..973L", "2005MNRAS.357..945G", "2005MNRAS.358...49M", "2005MNRAS.358.1083W", "2005MNRAS.362..799M", "2006A&A...455..315T", "2006ApJ...637..200Y", "2006MNRAS.365...46O", "2006MNRAS.371..703S", "2006hstc.conf..209G", "2007ASPC..364..315B", "2007ApJ...667..202L", "2007MNRAS.382..498C", "2008A&A...486..951G", "2008A&A...489..885M", "2008MNRAS.385.1998K", "2009A&A...501.1269K", "2009ApJS..185..289R", "2010A&A...516A..13P", "2010A&A...524A..11S", "2010MNRAS.404.1639V", "2010MNRAS.408.2495S", "2011A&A...525A..71W", "2011A&A...527A..91P", "2011A&A...536A.105V", "2011ASPC..448...91A", "2012ApJ...747...69C", "2012MNRAS.424..157V", "2013ApJ...763L..25C", "2013MNRAS.428.2949K" ]	[ "10.1051/0004-6361/201322505", "10.48550/arXiv.1403.7009" ]	1403	1403.7009_arXiv.txt	\label{} Spectral libraries play an important role in different fields of astrophysics. In particular they serve as a reference for the classification and automatic analysis of large stellar spectroscopic surveys and are fundamental ingredients for models of stellar populations used to study the evolution of galaxies. Much of what we know about the formation, evolution, and current state of galaxies comes through studies of their starlight. In distant galaxies, where the ability to study their stellar populations star-by-star is compromised by crowding and blending due to poor resolution, we must resort to studying their integrated light. This typically means comparing colors or spectra to models of simple or composite stellar population models \citep[e.g.,][]{IAPmodel,Buzzoni94, Worthey94,Leitherer96, % Pegase,Starburst99,BC03, Pegase-HR,Granada05,Maraston05,Vazdekis10, Conroy12a}. These comparisons give insight into a galaxy's evolution process: stellar population analysis can reveal the detailed chemical composition and star-formation history of a galaxy \citep[e.g.,][]{Gonzalez93, Davies93, Trager00,Thomas05,Yamada06,Koleva13,Conroy13}. Modern stellar population models consist of three primary ingredients \citep[e.g.,][]{Vazdekis10}: % stellar isochrones that represent the location in the luminosity--effective temperature plane (and, as a consequence, also surface gravity--effective temperature) of stars with different masses, the same initial chemical composition, and age; an initial mass function that populates the isochrones with a sensible number of stars; and a stellar library. A spectral library is a collection of stellar spectra that share similar wavelength coverage and spectral resolution. The spectra change as a function of effective temperature ($T _{\mathrm{eff}}$), surface gravity ($\log g$), and metallicity ([Fe/H]). To reproduce galaxy spectra as precisely as possible, one requires a comprehensive stellar spectral library that covers the entire desired parameter space of $T _{\mathrm{eff}}$, $\log g$, and [Fe/H]. Moreover, extended wavelength coverage is strongly desirable, because different stellar phases contribute their light in different bands. For instance, cool giants contribute more light than warmer faint giant stars in the near-infrared, while the situation is reversed in the optical \citep{Frogel88, Maraston05}. Asymptotic giant branch (AGB) stars dominate the light of intermediate-aged stellar populations in the near-infrared but are unimportant in the optical \citep{CB91, Worthey94, Maraston98}. Detecting their presence requires \emph{broad wavelength coverage} in both the target and model spectra. Stellar spectral libraries can be classified into empirical and theoretical libraries, depending on how libraries are obtained. Both theoretical and empirical libraries have improved in recent years. Widely-used theoretical libraries in stellar population models include those of \citet{Kurucz93,Coelho05,Martins05,Munari05, Coelho07,Gustafsson08}, % and \citet{Allard11}. Theoretical libraries have the advantage of (nearly) unlimited resolution and selectable abundance patterns, which include not only scaled-solar abundances but also non-solar patterns. Unfortunately, theoretical libraries suffer from systematic uncertainties, as they rely on model atmospheres and require a reliable list of atomic and molecular line wavelengths and opacities \citep{Coelho05}. Empirical stellar libraries, on the other hand, have the advantage of being drawn from real, observed stars and therefore do not suffer this limitation. However, they frequently have relatively low resolution (with a few exceptions; see below) and are typically unable to reproduce the indices measured in giant elliptical galaxies, because they are based on local stars with Milky Way disk abundance patterns \citep{Reynier89, Worthey92}. Table~\ref{tablib} lists several previous empirical stellar libraries and their principal features. In the optical, there are the (among others) Lick/IDS \citep{Worthey97}, MILES \citep{miles06}, ELODIE \citep{Prugniel01,Prugniel04,Prugniel07}, STELIB \citep{LeBorgne03}, NGSL \citep{Gregg06}, and the \citet{Pickles98} libraries. Building stellar libraries in the near-IR is a challenging task, but pioneering work has been done by \citet{Lancon92,LW2000} and \citet{LM2002} (LW2000, LM2002); \citet{MQ08}; and \citet[IRTF-SpeX]{Rayner09}. However, spectral libraries with extended wavelength coverage at moderate resolution are still largely missing. \begin{table} \caption{A selection of previous empirical stellar libraries.}\label{tablib} \begin{center} \begin{tabular}{lrrrl}\hline Library&Resolution&Spectral range&Number &Reference\\ &\multicolumn{1}{c}{R=$\lambda/\Delta\lambda$}& \multicolumn{1}{c}{(nm)}& \multicolumn{1}{c}{of stars}&\\\hline STELIB & 2000 & 320--930 & 249 & \citet{LeBorgne03}\\ ELODIE & 10000 & 390--680 & 1388 & \citet{Prugniel01,Prugniel04}\\ & & & & \citet{Prugniel07}\\ INDO-US & 5000 & 346--946 & 1237& \citet{Valdes04}\\ MILES & 2000 & 352--750 & 985 & \citet{miles06}\\ IRTF-SpeX & 2000 & 800--5000& 210 & \citet{Rayner09}\\ NGSL & 1000 & 167--1025& 374 & \citet{Gregg06}\\ UVES-POP & 80000 & 307--1030& 300 & \citet{Bagnulo03}\\ LW2000 & 1100 & 500--2500& 100 & \citet{LW2000}\\\hline \end{tabular} \end{center} \end{table} With the high-efficiency, broad-wavelength coverage and high/moderate resolution of the X-shooter spectrograph at ESO's VLT \citep{XShooter}, we can now fill the gap between high-resolution theoretical stellar libraries and low-resolution empirical stellar libraries. To this end, we developed the X-shooter Spectral Library (XSL, PI: Trager), which is a survey of $\gtrsim700$ stars that cover the entire Hertzsprung--Russell (HR) diagram, including both cool (M dwarfs, M giants, C stars, long-period variables, etc.) and hot stars (up to late O stars) with wavelength coverage from $300$--$2480\, \mathrm{nm}$. This includes the entire near-ultraviolet (NUV) to near-infrared (NIR) range at $R\sim8000$--11000. Here, we present the first two of six periods of XSL data (from ESO Periods 84 and 85). We concentrate on the NUV--optical data (3000--10200 \AA) from the UVB and VIS arms of X-shooter in this paper and leave the NIR arm data for a forthcoming paper.	We are building a new, moderate-resolution stellar spectral library, the X-shooter Spectral Library (XSL). The pilot survey\footnote{\url{http://xsl.u-strasbg.fr/}} (this work) contains 237 unique stars covering the spectral range $\lambda\lambda$ 3100 -- 10185 \AA\ at a resolution $R \sim 10000$. We have identified a number of issues with the X-shooter pipeline and presented our solutions. A telluric library is built for telluric correction of the XSL data using a PCA-based method. Flux and wavelength calibrations are carefully performed and are shown to be consistent with published spectra. the X-shooter Spectral Library is still under construction, and the final database will contain more than 700 stars. This library will provide a vital tool for extragalactic astronomers to extract even more information from their observations than previously possible and will provide stellar astronomers with a unique empirical panchromatic spectral library for further studies of a wide range of stellar types.	14	3	1403.7009	We present the first release of the X-shooter Spectral Library (XSL). This release contains 237 stars. The spectra in this release span a wavelength range of 3000-10 200 Å and have been observed at a resolving power of R ≡ λ/ Δλ ~ 10 000. The spectra were obtained at ESO's 8-m Very Large Telescope (VLT). The sample contains O-M, long-period variable, C and S stars. The spectra are flux-calibrated and telluric-corrected. We describe a new technique for the telluric correction. The wavelength coverage, spectral resolution, and spectral type of this library make it well suited to stellar population synthesis of galaxies and clusters, kinematical investigation of stellar systems, and the study of the physics of cool stars. <P />Full Table 3 and Table A.1 are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A117">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A117</A>	false	[ "cool stars", "R ≡ λ/ Δλ", "stellar population synthesis", "stellar systems", "kinematical investigation", "anonymous ftp", "XSL", "Spectral Library", "spectral type", "clusters", "spectral resolution", "VLT", "galaxies", "A117\">http://cdsarc.u", "A117</A", "237 stars", "the X-shooter Spectral Library", "O-M, long-period variable, C and S stars", "Table A.1", "CDS" ]	8.973154	9.860313	-1
482823	[ "Klein, Kristopher G.", "Howes, Gregory G.", "TenBarge, Jason M.", "Podesta, John J." ]	2014ApJ...785..138K	[ "Physical Interpretation of the Angle-dependent Magnetic Helicity Spectrum in the Solar Wind: The Nature of Turbulent Fluctuations near the Proton Gyroradius Scale" ]	60	[ "Department of Physics and Astronomy, University of Iowa, Iowa City, IA 52242, USA", "Department of Physics and Astronomy, University of Iowa, Iowa City, IA 52242, USA", "IREAP, University of Maryland, College Park, MD 20742, USA", "Center for Space Plasma Physics, Space Science Institute, Boulder, CO 80301, USA" ]	[ "2014ApJ...789..106H", "2014ApJ...790L..20K", "2014arXiv1404.2913H", "2015ApJ...802....1R", "2015ApJ...802....2R", "2015ApJ...805...24N", "2015ApJ...808..119I", "2015ApJ...813L..30H", "2015PhPl...22c2903K", "2015RSPTA.37340145H", "2015RSPTA.37340147G", "2015RSPTA.37340152O", "2016AIPC.1720d0013P", "2016ApJ...833..212M", "2016MNRAS.463L..79T", "2017ApJ...839..122K", "2017JGRA..12211835P", "2017JPlPh..83a7002H", "2017PhPl...24j2314A", "2018AnGeo..36...47R", "2018AnGeo..36..527R", "2018ApJ...855..121V", "2018ApJ...856...49W", "2018PhRvE..97a3204A", "2019ApJ...884L..53W", "2019ApJ...887...63I", "2019JPlPh..85c9003S", "2019LRSP...16....5V", "2019PhRvX...9c1037G", "2019arXiv190305710T", "2020ApJ...895...83Q", "2020ApJ...897L...3H", "2020ApJ...899...74B", "2020ApJ...905..142P", "2020ApJS..246...55D", "2020ApJS..246...66B", "2020GeoRL..4789720Z", "2020JPlPh..86d9002K", "2020arXiv201105672Z", "2021A&A...656A..21D", "2021ApJ...906..123Z", "2021ApJ...909L...7H", "2021ApJ...912..101W", "2021ApJ...913...80S", "2021ApJ...919...19A", "2021ApJ...922..240B", "2021JGRA..12627413Z", "2022ApJ...924...41V", "2022ApJ...924...92Z", "2022ApJ...936..128M", "2022Atmos..13..173P", "2022ExA....54..473V", "2022FrP....10.2167N", "2022PhRvL.129p5101B", "2022RAA....22a5009Z", "2022SpWea..2003200W", "2023A&A...677L..16R", "2023ApJ...953..161Z", "2023RAA....23i5027S", "2023SSRv..219...74K" ]	[ "astronomy", "physics" ]	7	[ "plasmas", "solar wind", "turbulence", "waves", "Physics - Plasma Physics", "Astrophysics - Solar and Stellar Astrophysics", "Physics - Space Physics" ]	[ "1938RSPSA.164..476T", "1982JGR....87.6011M", "1982PhRvL..48.1256M", "1984ApJ...285..109H", "1986JGR....9113366L", "1986JPlPh..35..431G", "1992JGR....97.3103G", "1992wapl.book.....S", "1994JGR....99.6011S", "1994JGR....9911519G", "1995AdSpR..15h.329L", "1995ApJ...438..763G", "1997JGR...10214661H", "1998ApJ...500..978Q", "1998ApJ...507L.181L", "1999JGR...104.2521L", "1999JGR...10422331L", "2000ApJ...539..273C", "2001ApJ...554.1175M", "2002GeoRL..29.1839K", "2003MNRAS.345..325C", "2005PhRvL..94u5002B", "2006GeoRL..33.9101H", "2006LRSP....3....1M", "2006PhRvL..96k5002B", "2007ApJ...655..269L", "2008ApJ...682.1070B", "2008ApJ...685..646C", "2008JGRA..113.1106H", "2008JGRA..113.5103H", "2008PhPl...15e5904H", "2008PhRvL.100f5004H", "2008PhRvL.101q5005H", "2009ApJ...698..986P", "2009ApJS..182..310S", "2009NPGeo..16..219H", "2009PhPl...16e2105V", "2009PhRvL.102b5003P", "2009PhRvL.102w1102S", "2009PhRvL.103g5006K", "2009PhRvL.103p5003A", "2009PhRvL.103u1101B", "2010ApJ...709L..49H", "2010ApJ...711L..79C", "2010MNRAS.407L..31W", "2010PhRvL.104y5002C", "2010PhRvL.105m1101S", "2011ApJ...731...85H", "2011ApJ...734...15P", "2011ApJ...737L..41C", "2011ApJ...738...40H", "2011ApJ...742...41P", "2011GeoRL..38.5101N", "2011PhRvL.107c5004H", "2012ApJ...745L...9S", "2012ApJ...749...86H", "2012ApJ...753L..19H", "2012ApJ...755..159K", "2012ApJ...758..120C", "2012PhPl...19e5901T", "2012PhRvL.108z1103O", "2012PhRvL.109s1101P", "2012SSRv..172...23M", "2012SSRv..172..325H", "2012SSRv..172..373M", "2013ApJ...769...58R", "2013ApJ...771L..27T", "2013ApJ...773..163V", "2013ApJ...774..139T", "2013PhPl...20g2302H", "2013PhPl...20g2303N", "2013PhRvL.110v5002C", "2013SoPh..286..529P" ]	[ "10.1088/0004-637X/785/2/138", "10.48550/arXiv.1403.2306" ]	1403	1403.2306_arXiv.txt	\label{sec:Intro} Turbulence is ubiquitous in space and astrophysical plasmas. In the heliosphere, it plays a critical role in plasma heating and the scattering of cosmic rays and energetic solar particles. Given the availability of detailed \emph{in situ} measurements from a broad range of heliophysics missions, the near-Earth solar wind is an ideal environment in which to study the fundamental nature of plasma turbulence. In the solar wind, the primary focus is the transport of energy from large scale turbulent motions to the small length scales at which the turbulence may be dissipated and the energy ultimately converted to plasma heat. Key questions at the forefront of solar wind turbulence investigations are: (i) what are the characteristics of the turbulent fluctuations?; (ii) what are the properties of the nonlinear interactions that drive the turbulent cascade of energy?; and (iii) what are the physical dissipation mechanisms that damp the turbulence and ultimately lead to plasma heating? The answers to each of these questions are interdependent. In particular, the nature of the electromagnetic and plasma fluctuations underlying the turbulence constrains the possible dissipation mechanisms. The definitive determination of the characteristics of the turbulent fluctuations, however, is made difficult due to the fact that the bulk of current \emph{in situ} solar wind observations are single point measurements. The Taylor hypothesis \citep{Taylor:1938} is typically invoked, because of the super-\Alfvenic velocity of the solar wind near 1 AU, to transform from the observed frequency of the advected turbulent fluctuations to a length scale of the fluctuations. For low frequency fluctuations with $\omega \lesssim \Omega_i$ and with the exception of a small number of studies of solar wind turbulence employing analyses based on multi-spacecraft measurements, we have no knowledge of the frequency of the fluctuations in the rest frame of the solar wind plasma. This limitation drives the quest to exploit alternative means to illuminate the character of the turbulent fluctuations in the solar wind. A wide range of characteristics, including polarizations, helicities, and other transport ratios have previously been used to decipher the dynamics of space plasmas \citep{Song:1994,Gary:1992, Lacombe:1995}. Here we employ the magnetic helicity as a sensitive probe for turbulent fluctuations in the solar wind. The \emph{fluctuating magnetic helicity}, defined as $H_m'\equiv \int d^3\V{r}\ \delta \V{A} \cdot \delta\V{B}$, where $\delta \V{A}$ and $\delta \V{B}$ are respectively the fluctuations of the vector potential and magnetic field, was first proposed as a useful metric for studying turbulent fluctuations in the solar wind by \cite{Matthaeus:1982b}. The \emph{reduced fluctuating magnetic helicity}, $H_m ^{r'}$ \citep{Matthaeus:1982a}, is related to $H_m'$ but is derivable from single-point spacecraft measurements and is a function of spacecraft-frame frequency $\omega_{SC}$, \citep{Howes:2010a}: \begin{equation} \begin{split} H_m^{r'}(\omega_{SC})=&\sum_{\V{k}} \frac{i[\V{B}_2(\V{k})\V{B}^_3(\V{k}) -\V{B}_2^(\V{k})\V{B}_3(\V{k})]}{\omega_{SC}/V_{SW}}\\ &\times \delta[\omega_{SC}- (\V{k}\cdot \V{V_{SW}} + \omega)], \end{split} \end{equation} where $\V{V}_{SW}$ is the solar wind velocity. Normalizing $H_m ^{r'}$ by the trace power $\|\V{B}(k_1)\|^2$ yields the \emph{normalized reduced fluctuating magnetic helicity}, $\sigma_m'=k_1 H_m ^{r'}(k_1)/\|\V{B}(k_1)\|^2$, where $k_1\equiv \omega_{SC}/V_{SW}$ is the projection of the wavevector along the direction of the solar wind. The normalized reduced fluctuating magnetic helicity is bound between $-1$ and $1$. For simplicity, throughout this paper the simplified term ``magnetic helicity" is used instead of the more cumbersome ``normalized reduced fluctuating magnetic helicity". Spacecraft measurements of $\sigma_m'$ in the near-Earth solar wind at low spacecraft-frame frequencies, $f \ll 1 \mbox{ Hz}$, typically give values that fluctuate about zero, which was originally interpreted as an admixture of waves with left-handed $(\sigma_m' \simeq -1)$ and right-handed $(\sigma_m' \simeq +1)$ magnetic helicities \citep{Matthaeus:1982a}. Based on the eigenfunctions of the linear Vlasov-Maxwell dispersion relation, \cite{Gary:1986} later showed that, at large-scales $k \rho_i \ll 1$, the linear waves have very small intrinsic magnetic helicity, $\sigma_m' \simeq 0$, which eliminated the need for an explanation in terms of a mixture of waves with left- and right-handed helicities. At higher spacecraft-frame frequencies, $f \sim 1 \mbox{ Hz}$, \emph{in situ} measurements of the magnetic field fluctuations produced an identifiably nonzero value of $\sigma_m'$ \citep{Goldstein:1994,Leamon:1998a}. This was initially interpreted as evidence of the damping of left-hand polarized ion cyclotron waves (ICWs) and the persistence of right-hand polarized whistler waves. Subsequently, it was shown that an anisotropic turbulent spectrum of \Alfven waves and kinetic \Alfven waves (KAWs) naturally reproduces both the low and high frequency measurements \citep{Howes:2010a}. As KAWs have preferentially perpendicular wavevectors, $k_\perp > k_\parallel$, and ICWs and whistler waves require sufficiently large parallel wavevectors, $k_\parallel d_i \gtrsim 1$, where $\perp$ and $\parallel$ refer to orientation with respect to the \emph{local} mean magnetic field $\V{B_0}$, the KAW model has the added benefit of being compatible with the predominately perpendicular cascade of energy expected from anisotropic magnetized turbulence theories \citep{Goldreich:1995,Boldyrev:2006}. Two recent studies \citep{He:2011a,Podesta:2011a}, analyzed the magnetic helicity of solar wind fluctuations as a function of the wave period $T$ in the spacecraft frame and the angle $\theta$ between the solar wind velocity $\V{V_{SW}}$ and $\V{B_0}$, and discovered two distinct signatures at angles perpendicular and parallel to $\V{B_0}$ at spacecraft-frame frequencies $f \sim 1$ Hz. These observations have been interpreted as two separate populations of fluctuations at kinetic scales with wavevectors oriented parallel and perpendicular to the local mean magnetic field \citep{Podesta:2011b,He:2012}. We here explore the properties of a model for solar wind turbulence that reproduces the observed behavior of the magnetic helicity. In particular, we show that the properties of the model are tightly constrained by spacecraft observations of the magnetic helicity plotted as a function of period and angle. This study employs the \emph{synthetic spacecraft data method} \citep{Klein:2012}, in which synthetic time series measurements of the magnetic field---measurements that may be directly compared to spacecraft observations---are generated by sampling along a trajectory through the model plasma volume. The turbulent fluctuations in this model are derived from a distribution of linear wave modes with random phases, where the particular distribution of wave power is guided by modern plasma turbulence theory, and where the physical properties of the wave modes are derived from linear kinetic plasma physics. The model's ability to reproduce the observed magnetic helicity structure is used to constrain the underlying parameters describing the distribution of turbulent wave power, shedding light on the nature of the turbulent fluctuations in the weakly collisional solar wind plasma. The remainder of the paper is organized as follows. Section~\ref{sec:mag_hel} provides background on the germane aspects of recent \textit{in situ} magnetic helicity measurements as well as a discussion of the underlying linear kinetic plasma theory. A description of and justification for the underlying turbulence model is given in Section \ref{sec:turbulence}. The general method used to generate the synthetic data and the particular results are found in Section \ref{sec:Synth}. Discussion of these results and a conclusion follow in Sections \ref{sec:discussion} and \ref{sec:summary} respectively.	\label{sec:summary} Motivated by recent novel measurements of the normalized reduced fluctuating magnetic helicity in the solar wind, we undertook in this paper to illuminate the nature of the underlying turbulence through careful comparison of \emph{in situ} and synthetic time series. The parallel and perpendicular signatures which are seen when solar wind observations of magnetic helicity are segregated by period and angle between the local mean magnetic field and solar wind velocity are replicated by synthetic time series derived from physically motivated models of turbulence. The hypothesis upon which the model is based is that the two signatures result from two distinct wavemode populations in wavevector space, namely a quasi-perpendicular collection of turbulent Alfv\'en/kinetic \Alfven fluctuations and a quasi-parallel collection of either whistler or ion cyclotron waves. In comparing the \emph{in situ} and synthetic magnetic helicity plots, we have been able to constrain the model's five free parameters, including the ratio of power contained in quasi-parallel versus quasi-perpendicular modes as well as the power ratio of sunward to anti-sunward modes for each of these two populations. From these comparisons, we conjecture that the parallel signature is not due to a cascade of energy to smaller parallel scales but is more likely due to a local injection of parallel energy from temperature anisotropy instabilities. These parallel modes may not be turbulent, existing along side and having little or no interaction with the cascade of energy to smaller perpendicular scales. Magnetic helicity alone is not sufficient to determine the nature of the parallel fluctuations, but we make several suggestions for a future determination of their nature. We discuss the effects of aliasing and show that the reduction of magnetic helicity amplitude to zero at small periods, which has been interpreted as evidence for an increasingly balanced turbulent cascade at smaller scales, is likely caused by aliasing and is not a physical effect. In conclusion, careful comparisons between magnetic helicity measurements and synthetic spacecraft data is capable of providing useful insight into the underlying nature of solar wind turbulence.	14	3	1403.2306	Motivated by recent observations of distinct parallel and perpendicular signatures in magnetic helicity measurements segregated by wave period and angle between the local magnetic field and the solar wind velocity, this paper undertakes a comparison of three intervals of Ulysses data with synthetic time series generated from a physically motivated turbulence model. From these comparisons, it is hypothesized that the observed signatures result from a perpendicular cascade of Alfvénic fluctuations and a local, non-turbulent population of ion-cyclotron or whistler waves generated by temperature anisotropy instabilities. By constraining the model's free parameters through comparison to in situ data, it is found that, on average, ~95% of the power near dissipative scales is contained in a perpendicular Alfvénic cascade and that the parallel fluctuations are propagating nearly unidirectionally. The effects of aliasing on magnetic helicity measurements are considered and shown to be significant near the Nyquist frequency.	false	[ "temperature anisotropy instabilities", "magnetic helicity measurements", "Alfvénic fluctuations", "whistler waves", "Ulysses data", "synthetic time series", "wave period", "situ data", "comparison", "dissipative scales", "Alfvénic", "distinct parallel and perpendicular signatures", "a perpendicular Alfvénic cascade", "the local magnetic field", "%", "Ulysses", "a perpendicular cascade", "Nyquist", "a local, non-turbulent population", "angle" ]	12.475888	14.507602	2
437497	[ "Girardi, M.", "Aguerri, J. A. L.", "De Grandi, S.", "D'Onghia, E.", "Barrena, R.", "Boschin, W.", "Méndez-Abreu, J.", "Sánchez-Janssen, R.", "Zarattini, S.", "Biviano, A.", "Castro-Rodriguez, N.", "Corsini, E. M.", "del Burgo, C.", "Iglesias-Páramo, J.", "Vilchez, J. M." ]	2014A&A...565A.115G	[ "Fossil group origins. III. The relation between optical and X-ray luminosities" ]	29	[ "Dipartimento di Fisica dell'Università degli Studi di Trieste - Sezione di Astronomia, via Tiepolo 11, 34143, Trieste, Italy ; INAF - Osservatorio Astronomico di Trieste, via Tiepolo 11, 34143, Trieste, Italy", "Instituto de Astrofísica de Canarias, C/Vía Láctea s/n, 38205 La Laguna, Tenerife, Canary Islands, Spain; Departamento de Astrofísica, Universidad de La Laguna, Av. del Astrofísico Franciso Sánchez s/n, 38205 La Laguna (Tenerife), Canary Islands, Spain", "INAF - Osservatorio Astronomico di Brera, via E. Bianchi 46, 23807, Merate (LC), Italy", "Astronomy Department, University of Wisconsin, 475 Charter St., Madison, WI, 53706, USA; Alfred P. Sloan Fellow", "Instituto de Astrofísica de Canarias, C/Vía Láctea s/n, 38205 La Laguna, Tenerife, Canary Islands, Spain; Departamento de Astrofísica, Universidad de La Laguna, Av. del Astrofísico Franciso Sánchez s/n, 38205 La Laguna (Tenerife), Canary Islands, Spain", "Fundación Galileo Galilei - INAF, Rambla José Ana Fernández Perez 7, 38712 Breña Baja, La Palma, Canary Islands, Spain", "Instituto de Astrofísica de Canarias, C/Vía Láctea s/n, 38205 La Laguna, Tenerife, Canary Islands, Spain; Departamento de Astrofísica, Universidad de La Laguna, Av. del Astrofísico Franciso Sánchez s/n, 38205 La Laguna (Tenerife), Canary Islands, Spain; School of Physics and Astronomy, University of St Andrews (SUPA), North Haugh, St Andrews, KY16 9SS, UK", "NRC Herzberg Institute of Astrophysics, 5071 West Saanich Road, Victoria, BC, V9E 2E7, Canada", "Instituto de Astrofísica de Canarias, C/Vía Láctea s/n, 38205 La Laguna, Tenerife, Canary Islands, Spain; Departamento de Astrofísica, Universidad de La Laguna, Av. del Astrofísico Franciso Sánchez s/n, 38205 La Laguna (Tenerife), Canary Islands, Spain", "INAF - Osservatorio Astronomico di Trieste, via Tiepolo 11, 34143, Trieste, Italy", "Instituto de Astrofísica de Canarias, C/Vía Láctea s/n, 38205 La Laguna, Tenerife, Canary Islands, Spain; Departamento de Astrofísica, Universidad de La Laguna, Av. del Astrofísico Franciso Sánchez s/n, 38205 La Laguna (Tenerife), Canary Islands, Spain", "Dipartimento di Fisica e Astronomia \"G. Galilei\", Università di Padova, vicolo dell'Osservatorio 3, 35122, Padova, Italy; INAF - Osservatorio Astronomico di Padova, vicolo dell'Osservatorio 5, 35122, Padova, Italy", "Instituto Nacional de Astrofísica, Óptica y Electrónica (INAOE), Aptdo. Postal 51 y 216, 72000 Puebla, Pue., Mexico", "Instituto de Astrofísica de Andalucia - C.S.I.C., 18008, Granada, Spain; Centro Astronómico Hispano Alemán, C/ Jesús Durbán Remón 2-2, 04004, Almería, Spain", "Instituto de Astrofísica de Andalucia - C.S.I.C., 18008, Granada, Spain" ]	[ "2014A&A...565A.116Z", "2014A&A...571A..49G", "2014MNRAS.442.1578R", "2015A&A...581A..16Z", "2015ApJ...806..268I", "2015MNRAS.454..161K", "2016A&A...586A..40K", "2016A&A...586A..63Z", "2016ApJ...821...29L", "2016ApJ...824..140R", "2016ApJ...826..146B", "2016ApJS..224...33B", "2017ApJ...834..164B", "2017ApJ...845...45K", "2017MNRAS.464.4593T", "2018A&A...609A..48A", "2018A&A...618A.172C", "2018A&A...620A...5A", "2018ApJ...853..129K", "2018MNRAS.474..866V", "2019A&A...631A..16Z", "2020A&A...639A..97A", "2020A&A...641A..95L", "2021A&A...655A.103Z", "2021Univ....7..132A", "2022A&A...668A..38Z", "2023A&A...671A..83G", "2023A&A...673A.100C", "2024arXiv240617012C" ]	[ "astronomy" ]	5	[ "galaxies: clusters: general", "X-rays: galaxies: clusters", "cosmology: observations", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1956nsbs.book.....S", "1968cgcg.bookR....Z", "1976ApJ...203..297S", "1979ats..book.....K", "1982ham..book.....L", "1987MNRAS.225..155F", "1989ApJS...70....1A", "1989MNRAS.237..799C", "1990ApJ...364..104I", "1992nrfa.book.....P", "1992scma.conf.....F", "1993MNRAS.265..706P", "1994Natur.369..462P", "1995A&A...297..610B", "1995PASP..107..945F", "1996AJ....111.1748F", "1996MNRAS.278..697P", "1997ApJ...479...90V", "1997MNRAS.290..119L", "1998AJ....116.3040G", "1998ApJ...504...27M", "1999A&A...349..389V", "1999ApJ...514..133M", "2000ApJ...530...62G", "2000ApJS..129..435B", "2000IAUC.7432....3V", "2001AJ....122.1104Y", "2001AJ....122.2267E", "2001ASPC..238..269L", "2002AJ....123..485S", "2003AJ....125.2064G", "2003MNRAS.343..627J", "2004A&A...423..449P", "2004A&A...425..367B", "2004ApJ...617..879L", "2005A&A...433..431P", "2005A&A...441..893A", "2005AJ....130..968M", "2005ApJ...630L.109D", "2006AJ....132..514C", "2006MNRAS.372L..68K", "2007A&A...469..363B", "2007AJ....134.1551S", "2007ApJ...660..239K", "2007MNRAS.377..595K", "2007MNRAS.382..433D", "2008MNRAS.386.2345V", "2009AJ....137.3942L", "2009ApJ...699..315Y", "2009ApJS..182..543A", "2009ApJS..183..197W", "2009MNRAS.395..255M", "2009MNRAS.396..900C", "2009MNRAS.398.1549R", "2010A&A...517A..52D", "2010ApJ...708.1376V", "2010ApJS..187..272W", "2010ApJS..191..254H", "2011A&A...526A.105Z", "2011A&A...527A.143A", "2011A&A...536A..89G", "2011MNRAS.416.2997C", "2011MNRAS.418.2054P", "2012A&A...537A..25M", "2012A&A...540A..90B", "2012ApJ...747...94M", "2012ApJ...752...12H", "2012MNRAS.422.3010L", "2014A&A...565A.116Z" ]	[ "10.1051/0004-6361/201323311", "10.48550/arXiv.1403.0590" ]	1403	1403.0590_arXiv.txt	\label{intro} Several studies of galaxy systems have revealed an interesting class of objects termed fossil groups (Ponman et al. \cite{pon94}). From the observational point of view, these are defined as galaxy systems with a magnitude difference of at least two magnitudes---in the $R$-band---between the brightest group/cluster galaxy (BCG) and the second-brightest galaxy within half the virial radius $R_{200}$\footnote{The radius $R_{\delta}$ is the radius of a sphere with mass overdensity $\delta$ times the critical density at the redshift of the galaxy system.} and an extended thermal X-ray halo with bolometric X-ray luminosity $L_{\rm X}{\rm (bol)}> 10^{42}\ h_{50 }^{-2}$ erg s $^{-1}$ (see Jones et al. \cite{jon03} for the rationale). Thus, the fossil groups appear to be extreme environments devoid of typical bright galaxies while simultaneously being home to the brightest and most massive galaxies in the Universe. The first explanation was that they are old, isolated galaxy systems in which the large galaxies have merged or coalesced through dynamical friction. In this merging scenario, the magnitude gap shown by the fossil systems is a consequence of evolution rather than an initial deficit of $\sim L^*$ galaxies (i.e., the failed group scenario; see, e.g., the discussion in the study of Mulchaey \& Zabludoff \cite{mul99}). The merging scenario has been invoked to explain such observational features as the high values of X-ray luminosity ($L_{\rm X}$) and temperature ($T_{\rm X}$) of fossil systems with respect to those of normal systems with comparable optical luminosity ($L_{\rm opt}$) or comparable velocity dispersion ($\sigma_{v}$; six fossil groups in Jones et al. \cite{jon03}; seven in Khosroshahi et al. \cite{kho07}) and some evidence of a high centrally concentrated dark matter halo (Khosroshahi et al. \cite{kho06}). The above differences with normal systems have been generally interpreted as due to an early formation epoch of fossil groups as suggested by numerical simulations (e.g., D'Onghia et al. \cite{don05}). Accordingly, the BCGs of fossil groups should contain a fossil relic of the structure formation in the high-redshift Universe. Early observations have revealed that the BCGs of fossil groups have different observational properties than other bright elliptical galaxies, their discy isophotes (seven fossil groups; Khosroshahi et al. \cite{kho06}), for example, supporting the idea that they are formed from gas-rich mergers in early times. More recent studies have opened the discussion about the special nature of fossil groups. Alternative criteria for their definition (e.g., Dariush et al. \cite{dar07}) and the concept of fossil clusters for massive systems (e.g., Cypriano et al. \cite{cyp06}) have been proposed. Moreover, studies based on $N$-body numerical simulations have suggested that many systems go through an optical fossil phase during their life (e.g., von Benda-Beckmann et al. \cite{von08}; Cui et al. \cite{cui11}). Recent observational results are often in contrast with the previous results that found no particularly high mass concentration (Democles et al. \cite{dem10}) and no special X-ray properties (12 fossil systems, Voevodkin et al. \cite{voe10}; 10, Proctor et al. \cite{pro11}; 17, Harrison et al. \cite{har12}). Instead, Proctor et al. (\cite{pro11}) claim atypical richnesses and optical luminosities, but this has not been found by Voevodkin et al. (\cite{voe10}) and Harrison et al. (\cite{har12}). Recent studies of fossil systems have also challenged the former conclusions of an early formation of their BCGs from a gas-rich merger. Analyzing the photometric and structural properties of BCGs in fossil systems, La Barbera et al. (\cite{lab09}, 25 fossil systems) and M\'endez-Abreu et al. (\cite{men12}, 20 fossil systems, hereafter Paper II) have found that they are similar to bright field ellipticals and to normal cluster BCGs, respectively. Finally, there is sparse evidence of a few fossil systems far from being dynamically relaxed (e.g., Harrison et al. \cite{har12}; La Barbera et al. \cite{lab12}; Miller et al. \cite{mil12}). Summarizing, there is still an open discussion on the real nature and origin of fossil systems. For instance, on the basis of their observational results, Harrison et al. (\cite{har12}) suggest that fossil systems formed rather early and their galaxies represent the end products of galaxy mergers, while Proctor et al. (\cite{pro11}) question the merging scenario, suggesting that the cannibalism of bright central galaxies is not a convincing explanation for the magnitude gap. Possible causes of the discrepancies among observational results reported in the literature might be connected with the use of very small samples, the presence of possible biases in the estimates of physical quantities, or inhomogeneities in the treatment of data of fossil and normal systems. In 2008 we started a large observational program of fossil systems, the FOssil Group Origins (FOGO) project (Aguerri et al. \cite{agu11}; hereafter Paper I). The aim of this project is to carry out a systematic, multiwavelength study of a sample of 34 fossil group candidates identified by Santos et al. (\cite{san07}, hereafter S07); here each system is denoted by FGS01, FGS02, etc., according to the S07 list. The FOGO project was awarded time as International Time Programme (ITP08-4 and ITP09-1) at the Roque de los Muchachos Observatory for a total of 52 nights of observations. Most optical and NIR observations were performed during the period November 2008--May 2010 at the TNG, NOT, WHT, and INT telescopes. The spectroscopic observations went on until April 2012 thanks to additional time awarded at TNG through the Spanish and Italian Time Allocation Committees. The catalog is described in the companion study by Zarattini et al. (\cite{zar13}; hereafter Paper IV). The first group we analyzed, RX J105453.3+552102 (FGS10 in the S07 catalog), is a special system, because it is already a very massive, relaxed galaxy cluster ($M\sim 1$ \mquii) at $z=0.47$. Contrary to the findings of previous works that claim a boost in the X-ray properties in fossil systems, FGS10 is quite normal as shown by its position in the $L_{\rm opt}$--$L_{\rm X}$ plane (see Paper I). Here we present our statistical results for 28 out of the 34 groups catalogued as fossils by S07. We have taken care to apply homogeneous procedures to the fossil and comparison systems and, in particular, we have computed consistent optical luminosities. Our present analysis is mainly based on optical data from the Sloan Digital Sky Survey Data Release 7 (hereafter SDSS-DR7, Fukugita et al. \cite{fuk96}; Gunn et al. \cite{gun98}; Abazajian et al. \cite{aba09}) and X-ray data from the {\em ROSAT} All Sky Survey (RASS, Voges et al. \cite{vog99}). We have also used the results of paper IV and, in particular, our check of the fossil classification of the S07 objects. This paper is organized as follows. We describe the S07 sample and the comparison sample in Sect.~2. We detail the computation of X-ray and optical luminosities in Sects.~3 and 4. We devote Sect.~5 to the comparison between fossil and normal systems in the $L_{\rm opt}$--$L_{\rm X}$ plane. We discuss our results and present our conclusions in Sect.~6. Unless otherwise stated, we indicate errors at the 68\% confidence level (hereafter c.l.). Throughout this paper, we use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ and $h_{70}=H_0/(70$ km s$^{-1}$ Mpc$^{-1}$) in a flat cosmology with $\Omega_{\rm m} =0.3$ and $\Omega_{\Lambda}=0.7$. Unless otherwise stated, all cosmology-dependent quantities that we take from the literature are rescaled to our adopted cosmology.		14	3	1403.0590	<BR /> Aims: This study is part of the Fossil group origins (FOGO) project which aims to carry out a systematic and multiwavelength study of a large sample of fossil systems. Here we focus on the relation between the optical luminosity (L<SUB>opt</SUB>) and X-ray luminosity (L<SUB>X</SUB>). <BR /> Methods: Out of a total sample of 28 candidate fossil systems, we consider a sample of 12 systems whose fossil classification has been confirmed by a companion study. They are compared with the complementary sample of 16 systems whose fossil nature has not been confirmed and with a subsample of 102 galaxy systems from the RASS-SDSS galaxy cluster survey. Fossil and normal systems span the same redshift range 0 <z< 0.5 and have the same L<SUB>X</SUB> distribution. For each fossil system, the L<SUB>X</SUB> in the 0.1-2.4 keV band is computed using data from the ROSAT All Sky Survey to be comparable to the estimates of the comparison sample. For each fossil and normal system we homogeneously compute L<SUB>opt</SUB> in the r-band within the characteristic cluster radius, using data from the Sloan Digital Sky Survey Data Release 7. <BR /> Results: We sample the L<SUB>X</SUB>-L<SUB>opt</SUB> relation over two orders of magnitude in L<SUB>X</SUB>. Our analysis shows that fossil systems are not statistically distinguishable from the normal systems through the 2D Kolmogorov-Smirnov test nor the fit of the L<SUB>X</SUB>-L<SUB>opt</SUB> relation. Thus, the optical luminosity of the galaxy system does strongly correlate with the X-ray luminosity of the hot gas component, independently of whether the system is fossil or not. We discuss our results in comparison with previous literature. <BR /> Conclusions: We conclude that our results are consistent with the classical merging scenario of the brightest galaxy formed via merger/cannibalism of other group galaxies with conservation of the optical light. We find no evidence for a peculiar state of the hot intracluster medium. <P />Tables 1 and 2 are available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201323311/olm">http://www.aanda.org</A>	false	[ "fossil systems", "systems", "other group galaxies", "X-ray luminosity", "Fossil", "28 candidate fossil systems", "opt</SUB", "each fossil system", "L<SUB", "102 galaxy systems", "the galaxy system", "previous literature", "the Sloan Digital Sky Survey Data Release", "data", "the normal systems", "L<SUB>X</SUB", "the X-ray luminosity", "each fossil and normal system", "comparison", "the comparison sample" ]	12.767176	5.411418	-1
2112714	[ "Bleem, L. E.", "Stalder, B.", "Brodwin, M.", "Busha, M. T.", "Gladders, M. D.", "High, F. W.", "Rest, A.", "Wechsler, R. H." ]	2015ApJS..216...20B	[ "A New Reduction of the Blanco Cosmology Survey: An Optically Selected Galaxy Cluster Catalog and a Public Release of Optical Data Products" ]	37	[ "Argonne National Laboratory, 9700 South Cass Avenue, Argonne, IL 60439, USA ; Kavli Institute for Cosmological Physics, University of Chicago, 5640 South Ellis Avenue, Chicago, IL 60637, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Department of Physics and Astronomy, University of Missouri, 5110 Rockhill Road, Kansas City, MO 64110, USA", "Kavli Institute for Particle Astrophysics and Cosmology 452 Lomita Mall, Stanford University, Stanford, CA 94305, USA ; SLAC National Accelerator Laboratory, 2575 Sand Hill Road, MS 29, Menlo Park, CA 94025, USA", "Kavli Institute for Cosmological Physics, University of Chicago, 5640 South Ellis Avenue, Chicago, IL 60637, USA ; Department of Astronomy and Astrophysics, University of Chicago, 5640 South Ellis Avenue, Chicago, IL 60637, USA", "Kavli Institute for Cosmological Physics, University of Chicago, 5640 South Ellis Avenue, Chicago, IL 60637, USA ; Department of Astronomy and Astrophysics, University of Chicago, 5640 South Ellis Avenue, Chicago, IL 60637, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Kavli Institute for Particle Astrophysics and Cosmology 452 Lomita Mall, Stanford University, Stanford, CA 94305, USA ; SLAC National Accelerator Laboratory, 2575 Sand Hill Road, MS 29, Menlo Park, CA 94025, USA ; Department of Physics, Stanford University, Stanford, CA 94305, USA" ]	[ "2014ApJ...797..109R", "2015ApJS..216...27B", "2015MNRAS.453..311S", "2015MNRAS.454.2305S", "2016A&A...592A..11K", "2016ApJ...830..143W", "2016MNRAS.456.3213G", "2016MNRAS.460.3900F", "2016MNRAS.462.4141L", "2016arXiv160607039D", "2017ApJ...836....9C", "2017MNRAS.464.2910S", "2017MNRAS.468..577S", "2017NewA...51..169S", "2018A&A...614A..82T", "2018A&A...620A..20K", "2018ApJ...867...12R", "2018MNRAS.473.1258S", "2018MNRAS.475..343W", "2018MNRAS.481.5247O", "2018PASJ...70S..20O", "2019ApJ...872..170R", "2019ApJS..243...17J", "2019PhRvD.100d3501O", "2020AJ....159..110H", "2020ApJS..247...25B", "2020MNRAS.491.5852A", "2020MNRAS.495..705Z", "2021A&A...647A.106M", "2021ApJ...910...60S", "2021ApJS..253....3H", "2022A&A...665A..78S", "2022ApJS..258...36B", "2022SSPMA..52B9503H", "2023A&A...673A..92B", "2023ApJ...947...44C", "2024PASA...41...22S" ]	[ "astronomy" ]	6	[ "galaxies: clusters: general", "surveys", "techniques: photometric", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1955ApJ...121..161S", "1958ApJS....3..211A", "1972CoASP...4..173S", "1989ApJS...70....1A", "1994A&AS..104..365F", "1996A&AS..117..393B", "1998ApJ...502..558V", "2000AJ....120.1579Y", "2000AJ....120.2148G", "2002ASPC..281..228B", "2003MNRAS.344.1000B", "2004A&A...425..367B", "2004cgpc.sympE..53W", "2005ApJ...634.1103R", "2005ApJS..157....1G", "2006AJ....131.1163S", "2007AJ....134..973I", "2007AJ....134.2398C", "2007ApJ...654..858B", "2007ApJ...660..221K", "2007ApJ...660..239K", "2007ApJ...666..674M", "2008ApJ...676..868D", "2009AJ....138..110H", "2009ApJ...690...42M", "2009ApJ...698.1221M", "2009ApJ...699..768R", "2009ApJ...701..428A", "2009ApJS..183..197W", "2009MNRAS.393.1275G", "2009PhRvD..79f3009C", "2010AIPC.1279..120T", "2010ApJ...713.1207W", "2010ApJ...715..823G", "2010ApJ...722.1180V", "2010ApJ...723.1736H", "2010ApJS..191..254H", "2010ApJS..191..340M", "2010SPIE.7733E..0EK", "2011A&A...536A..12P", "2011AAS...21724907M", "2011AJ....141...94G", "2011ARA&A..49..409A", "2011ApJ...727...45S", "2011ApJ...736...21S", "2011ApJ...737...61M", "2011ApJ...742...48B", "2011MSAIS..17..159C", "2011PASP..123..568C", "2012A&A...537A..39S", "2012AJ....143...38G", "2012ApJ...746..178R", "2012ApJ...753L...9B", "2012ApJ...757...83D", "2012ApJ...761...22S", "2012ApJS..199...34W", "2012ApJS..200....9G", "2012MNRAS.423..909C", "2012SPIE.8452E..1EA", "2013ApJ...762...10D", "2013ApJ...763..127R", "2013ApJ...767...38S", "2013ApJ...771L..16H", "2013ApJS..209...22A", "2014A&A...571A..29P", "2014ApJ...785..104R", "2014ApJ...792...45R", "2014MNRAS.438...78R" ]	[ "10.1088/0067-0049/216/1/20", "10.48550/arXiv.1403.7186" ]	1403	1403.7186_arXiv.txt	\setcounter{footnote}{0} Multi-band optical surveys provide powerful data sets with which to study cosmology. Technological improvements in CCD imagers over the past two decades have led to a rapid increase in the number of such surveys. Notable examples include the large-area Sloan Digital Sky Survey \citep[SDSS][]{york2000} and deeper, moderate-area surveys including the Red Sequence Cluster Survey \citep{gladders05}, the Canada-France-Hawaii Telescope Legacy Survey\footnote{http://www.cfht.hawaii.edu/Science/CFHLS/}, and the NOAO Deep Wide-Field Survey \citep{jannuzi99}. A new generation of both large ($\gtrsim 1000$ \degs) and deep surveys including the Panoramic Survey Telescope \& Rapid Response System \citep{kaiser10}, RCS2 \citep{gilbank11}, the Dark Energy Survey\footnote{www.darkenergysurvey.org} and the Subaru Hyper Suprime-Cam project \citep{takada10} highlight further advances in the field. Such surveys are particularly useful for constraining cosmology with galaxy clusters as large areas are required to obtain statistically useful numbers of these rare, massive objects. The selection of galaxy clusters from optical surveys has a rich history beginning with Abell's visual identification of over-densities of galaxies in the Palomar Sky Survey \citep{abell58} and continuing with modern techniques that identify concentrations of the red, passively-evolving E/S0 galaxies that form the red-sequence in clusters (e.g., \citealt{gladders00, koester07a, hao10}; see \citet{allen11} for a recent review). The critical challenge for cosmology with clusters is connecting observable cluster properties to the mass of the system. While optical selection does provide a mass proxy --- usually in the form of a ``richness'' parameter related to the number of galaxies in the cluster --- complimentary mm-wave or X-ray data sets can greatly improve cluster mass-calibration and provide a cross-check by testing the self-consistency of the respective observable-to-mass scaling relations \citep{rozo12c}. Indeed, joint optical and mm-wave analyses of the maxBCG cluster sample \citep{koester07a} with data from the \emph{Planck} satellite \citep{planck11-5.2c} and the Atacama Cosmology Telescope \citep[ACT][]{sehgal12} have already revealed tensions between the mm-wave mass and optical richness-mass scaling relations. The Blanco Cosmology Survey \citep[BCS][]{desai12}, an \bcssize \ multi-band optical survey, while of moderate depth, is particularly interesting owing to the wealth of multi-wavelength data in the survey region. In addition to being contained within the footprint of both the South Pole Telescope \citep[SPT][]{carlstrom11} and ACT mm-wave surveys, one of the two BCS fields contains the XMM-BCS survey \citep{suhada12}, a subset of a new \emph{XMM} survey, the XXL\footnote{http://irfu.cea.fr/xxl/}, and is itself contained within the 100 \degs \ SPT Deep Field which features near-infrared imaging from \emph{Spitzer} \citep{ashby13} and has been imaged with \emph{Herschel}'s SPIRE camera at 250, 350 and 500 microns \citep{holder13}. In this paper we present calibrated optical source catalogs from the BCS. We identify galaxy clusters in these data using a cluster-detection technique based upon the methodology of \citet{gladders00}, which uses a matched filter in position, color and magnitude space to identify spatial over-densities of red-sequence galaxies. The paper is organized as follows: in \S \ref{sec:reduction}, we provide a brief overview of the BCS and describe our image reduction pipeline and photometric calibration. We describe the galaxy cluster detection algorithm in \S \ref{sec:cluster}, and measurement of cluster properties in \S \ref{sec:remeasure}. In \S \ref{sec:tests} we discuss tests of our algorithm on simulated catalogs and in \S \ref{sec:compare} compare our results to previous cluster catalogs that overlap the BCS region. Finally, we conclude in \S \ref{sec:discussion}. Where applicable we assume a flat $\Lambda$CDM cosmology with $\Omega_{M}=0.27$ and $h=0.71$. Unless otherwise specified, all masses are reported in terms of \mtwohundred, where \mtwohundred \ is defined as the mass contained within a radius $r_{200}$ at which the average density is 200 times the critical density, and all magnitudes are reported in the AB system.	\label{sec:discussion} In this paper we have presented our reductions of optical imaging data from the Blanco Cosmology Survey, including our methodology for creating calibrated source catalogs, for correcting underestimated photometric errors returned by SExtractor, and our implementation of a new, easily-coded morphological star-galaxy separation statistic. Using a red-sequence-based cluster finding algorithm, we have searched the BCS for galaxy clusters. We report the coordinates, redshifts, and optical richnesses for \numberclusters \ clusters at z $\leq$\redshiftlimit, of which greater than \newpercent \% are new detections. This sample has a median redshift of \medianredshift \ and median optical richness, $\lambda$, of \medianrichness. Based upon tests with realistic mock catalogs, the catalog is expected to be $>85\%$ pure at z $< 0.75$. One of the strengths of the Blanco Cosmology Survey is its overlap with other multi-wavelength data. The optical catalogs presented here have been used to confirm cluster candidates and estimate redshifts for SZ-selected clusters in the SPT-SZ survey \citep{reichardt12, song12b}, as well as to make the first measurement of galaxy bias from the gravitational lensing of the Cosmic Microwave Background \citep{bleem12b}. As these reductions might be of broader utility, we release the reduced \emph{g-}, \emph{r-}, \emph{i-}, \emph{z-}band images and weight maps as well as the calibrated source catalogs for this \bcssize \ survey. These products are available at \webaddress.	14	3	1403.7186	The Blanco Cosmology Survey is a four-band (griz) optical-imaging survey of ~80 deg<SUP>2</SUP> of the southern sky. The survey consists of two fields centered approximately at (R.A., decl.) = (23<SUP>h</SUP>, -55°) and (5<SUP>h</SUP>30<SUP>m</SUP>, -53°) with imaging sufficient for the detection of L <SUB>sstarf</SUB> galaxies at redshift z <= 1. In this paper, we present our reduction of the survey data and describe a new technique for the separation of stars and galaxies. We search the calibrated source catalogs for galaxy clusters at z <= 0.75 by identifying spatial over-densities of red-sequence galaxies and report the coordinates, redshifts, and optical richnesses, λ, for 764 galaxy clusters at z <= 0.75. This sample, >85% of which are new discoveries, has a median redshift of z = 0.52 and median richness λ(0.4 L <SUB>sstarf</SUB>) = 16.4. Accompanying this paper we also release full survey data products including reduced images and calibrated source catalogs. These products are available at <A href="http://data.rcc.uchicago.edu/dataset/blanco-cosmology-survey">http://data.rcc.uchicago.edu/dataset/blanco-cosmology-survey</A>.	false	[ "galaxy clusters", "redshift z", "galaxies", "lt;=", "optical richnesses", "z", "redshifts", "median richness", "redshift z <=", "red-sequence galaxies", "spatial over-densities", "z <=", "764 galaxy clusters", "calibrated source catalogs", "full survey data products", "a median redshift", "L <SUB>sstarf</SUB> galaxies", "new discoveries", "decl", "catalogs" ]	12.638788	5.127037	-1
437290	[ "Ganse, U.", "Kilian, P.", "Spanier, F.", "Vainio, R." ]	2014A&A...564A..15G	[ "Fundamental and harmonic plasma emission in different plasma environments" ]	11	[ "Department of PhysicsUniversity of Helsinki, 00014, Helsinki, Finland", "Max-Planck Institut für Sonnensystemforschung, 37077, Göttingen, Germany", "Center for Space Research, North-West University, 2520, Potchefstrom, South Africa", "Department of PhysicsUniversity of Helsinki, 00014, Helsinki, Finland; Department of Physics and Astronomy, University of Turku, 20014, Turku, Finland" ]	[ "2014PhPl...21h3109T", "2015A&A...584A..83T", "2015ApJ...798...81K", "2016A&A...586A..19T", "2018CoPhC.224..245M", "2018PhPl...25k3110A", "2020ApJ...904...88A", "2021Ap&SS.366...60Z", "2021ApJ...913...99K", "2021PPCF...63d5001V", "2022arXiv220700462J" ]	[ "astronomy" ]	6	[ "Sun: radio radiation", "waves", "plasmas", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1950AuSRA...3..387W", "1970AuJPh..23..871M", "1981PlPh...23..413M", "1986islp.book.....M", "1988csup.book.....H", "2000PhPl....7.3167W", "2001JGR...10625041K", "2002A&A...384..273A", "2003JGRA..108.1361K", "2005ApJ...623.1180C", "2006SSRv..123..251F", "2008ApJ...674.1078O", "2008ApJ...676.1330P", "2009ApJ...691L.151L", "2012ApJ...751..145G", "2012SoPh..280..551G", "2013A&A...558A...7P" ]	[ "10.1051/0004-6361/201322834", "10.48550/arXiv.1403.2240" ]	1403	1403.2240_arXiv.txt			14	3	1403.2240	<BR /> Aims: Emission of radio waves from plasmas through plasma emission with fundamental and harmonic frequencies is a familiar process known from solar type II radio bursts. Current models assume the existence of counterstreaming electron beam populations excited at shocks as sources for these emission features, which limits the plasma parameters to reasonable heliospheric shock conditions. However, situations in which counterstreaming electron beams are present can also occur with different plasma parameters, such as higher magnetisation, including but not limited to our Sun. Similar radio emissions might also occur from these situations. <BR /> Methods: We used particle-in-cell simulations to compare plasma microphysics of radio emission processes from counterstreaming beams in different plasma environments that differed in density and magnetization. <BR /> Results: Although large differences in wave populations are evident, the emission process of type II bursts appears to be qualitatively unaffected and shows the same behaviour in all environments.	false	[ "different plasma environments", "plasma emission", "radio emission processes", "different plasma parameters", "plasma microphysics", "solar type II radio bursts", "Similar radio emissions", "plasmas", "magnetization", "reasonable heliospheric shock conditions", "Emission", "electron beam populations", "radio waves", "electron beams", "density", "shocks", "higher magnetisation", "beams", "wave populations", "Sun" ]	13.522057	15.258112	2
812905	[ "Yagi, Kent", "Kyutoku, Koutarou", "Pappas, George", "Yunes, Nicolás", "Apostolatos, Theocharis A." ]	2014PhRvD..89l4013Y	[ "Effective no-hair relations for neutron stars and quark stars: Relativistic results" ]	113	[ "Department of Physics, Montana State University, Bozeman, Montana 59717, USA", "Department of Physics, University of Wisconsin-Milwaukee, P.O. Box 413, Milwaukee, Wisconsin 53201, USA", "School of Mathematical Sciences, The University of Nottingham, University Park, Nottingham NG7 2GD, United Kingdom; SISSA, Via Bonomea 265, 34136 Trieste, Italy", "Department of Physics, Montana State University, Bozeman, Montana 59717, USA", "Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis Zografos GR15783, Athens, Greece" ]	[ "2014ApJ...788...15S", "2014ApJ...791...78A", "2014PhRvD..90b4025P", "2014PhRvD..90f1501K", "2014PhRvD..90f3010Y", "2014PhRvD..90f4026M", "2014PhRvD..90f4030C", "2014PhRvD..90j4021D", "2014PhRvD..90l4023C", "2014arXiv1409.2472B", "2015AIPC.1693c0001B", "2015ApJ...798..121S", "2015ApJ...798L..17C", "2015CQGra..32n5008S", "2015CQGra..32x3001B", "2015JHEP...06..059L", "2015JHEP...09..219L", "2015MNRAS.453.2862P", "2015MNRAS.454.4066P", "2015PhRvD..91d4011P", "2015PhRvD..91d4017C", "2015PhRvD..91f4001T", "2015PhRvD..91j3003Y", "2015PhRvD..91l3008Y", "2015PhRvD..92b3007C", "2015PhRvD..92b4010P", "2015PhRvD..92b4020M", "2015PhRvD..92b4056G", "2015PhRvD..92f4015D", "2015PhRvD..92h3009B", "2015PhRvD..92j4008C", "2015PhRvD..92l4001D", "2015PhRvD..92l4003P", "2015PhRvD..92l4030P", "2015arXiv150407335D", "2015arXiv151206768S", "2016ApJ...818L..11T", "2016CQGra..33i5005Y", "2016CQGra..33q4001C", "2016EJPh...37f5602B", "2016GrCo...22..305B", "2016JCAP...01..008L", "2016MNRAS.459..646B", "2016MNRAS.459.4378S", "2016PhRvD..93b4004S", "2016PhRvD..93b4033C", "2016PhRvD..93f4077K", "2016PhRvD..93j4051M", "2016PhRvD..93l4041M", "2016PhRvD..94f4015U", "2016arXiv161108814M", "2017CQGra..34a5006Y", "2017CQGra..34q7002M", "2017LRR....20....7P", "2017MNRAS.466.4381P", "2017PhR...681....1Y", "2017PhRvD..96b3005M", "2017RPPh...80i6901B", "2018ASSL..457..575H", "2018ASSL..457..737D", "2018ApJ...861..141L", "2018CQGra..35a5005S", "2018CQGra..35b5009G", "2018MNRAS.478.1377A", "2018MNRAS.478.1893B", "2018PhRvD..97b4043G", "2018PhRvD..97d1502G", "2018RSOS....580640F", "2018arXiv180603904M", "2019ApJ...877...66U", "2019CQGra..36a5010A", "2019CQGra..36n7002M", "2019EPJP..134..600E", "2019GReGr..51...46F", "2019PhRvD..99b4041G", "2019PhRvD..99d4007N", "2019PhRvD..99j4014P", "2019PhRvD..99j4050R", "2019PhRvD..99l3026K", "2019PhRvD.100b1501S", "2019PhRvD.100d4003D", "2019PhRvD.100l4021M", "2019PrPNP.10903714B", "2019Symm...11.1324B", "2020ApJ...899..139M", "2020MNRAS.498.3616A", "2020NatCo..11.2553P", "2020PhRvD.101f4003B", "2020PhRvD.101l4006J", "2021GReGr..53...27D", "2021PhRvD.103b4049D", "2021PhRvD.103f3038S", "2021PhRvD.104f4012F", "2021arXiv210512582S", "2022CQGra..39m4001K", "2022PhRvD.105b3018T", "2022Univ....8..353A", "2023AN....34420109K", "2023ApJ...954...16G", "2023GReGr..55..123D", "2023PhLB..84037869M", "2023PhRvC.108b5810L", "2023PhRvD.107f4001G", "2023PhRvD.107j3050P", "2023PhRvD.108d3015A", "2023PhRvD.108d4042M", "2023arXiv231115277K", "2024EPJP..139..273B", "2024PhRvD.109b3020L", "2024PhRvD.109b4062A", "2024PhRvD.109j3029S", "2024RvMP...96a5004D", "2024arXiv240607613A" ]	[ "astronomy", "physics" ]	9	[ "04.40.Dg", "97.60.Jd", "Relativistic stars: structure stability and oscillations", "Neutron stars", "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena", "High Energy Physics - Phenomenology" ]	[ "1909RSPSA..82...73L", "1959PhRv..113..934K", "1967ApJ...150.1005H", "1967PhRv..164.1776I", "1968ApJ...153..807H", "1968CMaPh...8..245I", "1968PhRv..167.1175E", "1968PhRv..168.1415E", "1970JMP....11.2580G", "1971PhRvL..26..331C", "1971PhRvL..26.1344H", "1972CMaPh..25..152H", "1974JMP....15...46H", "1975PhRvL..34..905R", "1976ApJ...204..200B", "1977ApJS...33..415A", "1980RvMP...52..299T", "1984PhRvD..30.2379F", "1988PhRvC..38.1010W", "1989JMP....30.2252F", "1989MNRAS.237..355K", "1992ApJ...398..203C", "1993A&A...278..421B", "1993ApJS...88..205L", "1994A&A...291..155S", "1995ApJ...444..306S", "1995JMP....36.3063M", "1995PhRvD..52..821K", "1995PhRvD..52.5707R", "1997PhRvD..56.1845R", "1998A&AS..132..431N", "1998NuPhA.637..435S", "1998PThPh.100.1013S", "1998PhRvC..58.1804A", "1998PhRvD..58j4020B", "2000A&A...359..311Z", "2000PhRvD..61h1501M", "2001A&A...380..151D", "2001ApJ...550..426L", "2002MNRAS.336..831S", "2002Natur.420...51C", "2003LRR.....6....3S", "2004MNRAS.350.1416B", "2005CQGra..22.1607B", "2005MNRAS.358..923B", "2005PhRvD..72d4028B", "2006Natur.441.1115O", "2006Sci...311.1901H", "2007ApJ...663.1244M", "2008ApJ...677.1216H", "2008CoPhC.179..586M", "2008PhRvD..77b1502F", "2009GReGr..41.2415B", "2009PhRvD..80h4018B", "2009PhRvD..80h4035D", "2010PhRvD..81l3016H", "2010arXiv1003.5015G", "2012ApJ...753..175B", "2012MNRAS.422.2581P", "2012PhRvL.108w1104P", "2012SPIE.8443E..13G", "2012SPIE.8443E..2DF", "2013ApJ...766...87B", "2013ApJ...771...51L", "2013ApJ...777...68B", "2013MNRAS.429.3007P", "2013PhRvD..88b3009Y", "2013Sci...340..448A", "2013Sci...341..365Y", "2014ApJ...781L...6D", "2014ApJ...787..136P", "2014ApJ...788...15S", "2014ApJ...792...87P", "2014MNRAS.438L..71H", "2014PhRvD..89d3011Y", "2014PhRvL.112l1101P", "2014PhRvL.112t1102C" ]	[ "10.1103/PhysRevD.89.124013", "10.48550/arXiv.1403.6243" ]	1403	1403.6243_arXiv.txt	Stationary and axisymmetric black holes (BHs) satisfy certain no-hair relations~\cite{MTW,robinson,israel,israel2,hawking-uniqueness0,hawking-uniqueness,carter-uniqueness}, which allow us to completely describe them in terms of their mass, spin angular momentum and charge. Such relations imply that the exterior gravitational field of BHs can be expressed as an infinite series of multipole moments that only depend on the first two (in the absence of charge): the mass-monopole (the mass) and the current-dipole (the spin angular momentum)~\cite{geroch,hansen}. The multipole moments used to describe the gravitational potential far from a source are analogous to those used in electromagnetism to describe the electric and vector potential of a distribution of charge and current. As in electromagnetism, in General Relativity (GR) these multipoles come in two flavors: mass moments and current moments~\cite{thorne-MM}. The former is sourced by the energy density (or the time-time component and the trace of the spatial part of the matter stress-energy tensor) while the latter is sourced by the energy current density (or the time-space component of the stress-energy tensor)~\cite{kidder}. Multipole moments are important not only because they allow us to describe the exterior gravitational field~\cite{Backdahl:2005uz,Backdahl:2006ed}, but also because they are directly related to astrophysical observables~\cite{Ryan:1995wh,Ryan:1997hg,Pappas:2012nt}. The BH no-hair relations do not apply to neutron stars (NSs) and quark stars (QSs) since these are not vacuum solutions to the Einstein equations. One may then expect that the multipole moments of NSs and QSs would depend strongly on their internal structure, or more precisely, on their equation of state (EoS), which relates their internal pressure to the energy density. The goal of this paper is to investigate whether such NSs and QSs satisfy approximate no-hair relations in full GR. If this were the case, one would then be able to describe their exterior gravitational field by measuring their first few multipole moments, which could have interesting applications to X-ray NS observations~\cite{Baubock:2013gna,Psaltis:2013fha}. Some evidence already exists to support the existence of such no-hair relations for NSs and QSs. A universal relation between the moment of inertia (directly related to the current dipole moment) and the mass quadrupole moment was found in~\cite{I-Love-Q-Science,I-Love-Q-PRD}, using an unmagnetized, uniform- and slow-rotation approximation. This result was immediately confirmed by~\cite{lattimer-lim} using different EoSs. Haskell \textit{et al.}~\cite{I-Love-Q-B} extended the analysis of~\cite{I-Love-Q-Science,I-Love-Q-PRD} to magnetized NSs and found that the universality still holds, provided that stars spin moderately fast (spin period less than 0.1s) and the magnetic fields are not too large (less than $10^{12}$G). Such properties are precisely those one expects millisecond pulsars to have. Several studies have relaxed the slow-rotation approximation~\cite{doneva-rapid,Pappas:2013naa,Chakrabarti:2013tca}, leading to an apparent initial disagreement on whether rapid rotation can destroy the EoS-universality \emph{in general}. Doneva \textit{et al.}~\cite{doneva-rapid} constructed NS and QS sequences by varying the dimensional spin frequency and found, as anticipated in~\cite{I-Love-Q-Science,I-Love-Q-PRD}, that the relation between the moment of inertia and the quadrupole moment depend on the spin frequency. By constructing stellar sequences by varying the dimensional spin-frequency, they found that EoS-universality breaks down and from that, they concluded that \emph{in general} the EoS-universality is broken for rapidly rotating stars. However, shortly after, Pappas and Apostolatos~\cite{Pappas:2013naa} and Chakrabarti \textit{et al.}~\cite{Chakrabarti:2013tca} constructed stellar sequences by varying dimensionless combinations of the spin angular momentum and found that the relations remain EoS-universal. Although the conclusions of~\cite{doneva-rapid} and~\cite{Pappas:2013naa,Chakrabarti:2013tca} are contradictory, their calculations are not actually in disagreement; instead, they show that whether the EoS-universality holds for rapidly-rotating relativistic stars depends on the choice of spin parametrization. More recently, Stein \textit{et al.}~\cite{Stein:2013ofa} proved analytically that universality is preserved regardless of rotation, provided one works with dimensionless spin parameterizations. Recent studies have also considered whether approximately EoS independent relations exist between higher-$\ell$ multipole moments. Reference~\cite{Pappas:2013naa} in fact found one such relation between the current octupole and the mass quadrupole moments of NSs. This relation was not only approximately EoS-universal but also approximately spin-insensitive. The Newtonian results of~\cite{Stein:2013ofa} analytically confirmed this result. The latter, in fact, proved that higher-$\ell$ multipole moments in the non-relativistic Newtonian limit can be expressed in terms of just the mass monopole, spin current dipole and mass quadrupole moments through relations that are approximately EoS-universal and spin-independent. This universality, however, was found to deteriorate with increasing $\ell$ multipole number. The existence of approximately universal relations is not only of academic interest, but it also has practical applications. For example, if one could measure any two quantities in a given relation independently, one could perform an EoS-independent test of GR in the strong-field regime~\cite{I-Love-Q-Science,I-Love-Q-PRD}. Moreover, these relations may play a critical role when attempting to measure the mass and radius of NSs with future X-ray telescopes, such as NICER~\cite{2012SPIE.8443E..13G} and LOFT~\cite{2012AAS...21924906R,2012SPIE.8443E..2DF}. The pulse and atomic line profiles of such stars depend not only on the stellar mass and radius, but also on the moment of inertia, the quadrupole moment and the stellar eccentricity~\cite{Morsink:2007tv,Baubock:2012bj,Psaltis:2013zja}. Universal relations between these quantities~\cite{I-Love-Q-Science,I-Love-Q-PRD,Baubock:2013gna} allow one to break parameter degeneracies and measure the mass and radius~\cite{Psaltis:2013fha}. Such measurements, in turn, would allow for exquisite constraints on the EoS in the high density regime~\cite{2006Natur.441.1115O}. In this paper, we study whether approximately EoS-independent relations among multipole moments exist up to hexadecapole order in full GR for both NSs and QSs. To do so, we construct unmagnetized, uniformly-rotating NS and QS solutions to the Einstein equations. For rapidly-rotating stars, we extract multipole moments by numerically constructing stellar solutions with the LORENE~\cite{bonazzola_gsm1993,bonazzola_gm1998} and RNS~\cite{stergioulas_friedman1995} codes. For slowly-rotating stars, we extract multipole moments by solving the Einstein equations in a slow-rotation expansion to quartic order in spin, extending previously-found quadratic~\cite{hartle1967,Hartle:1968ht} and cubic~\cite{benhar} solutions. Validity of the quadratic solution is discussed in~\cite{berti-white}. Such an extension allows us to estimate the importance of quartic-order-in-spin terms in X-ray observations of millisecond pulsars, which were neglected in~\cite{Psaltis:2013zja,Psaltis:2013fha}. \begin{figure}[htb] \centering \includegraphics[width=\columnwidth,clip=true]{fast-rot-M4-M2.pdf} \caption{ (Color online) (Top) The (reduced dimensionless) hexadecapole $(\bar{M}_4)$--quadrupole $(\bar{M}_2)$ moments relation with various NS and QS EoSs and spins, together with the fit given by Eq.~\eqref{eq:fit} and the Newtonian relation found in~\cite{Stein:2013ofa}. Observe the relation approaches the Newtonian one as one increases $\bar{M}_2$. The Newtonian relation for an $n=0.5$ polytrope agrees with the relativistic fit for various realistic EoSs within 10\% accuracy above the critical $\bar{M}_2$ denoted by the dotted-dashed, vertical line. (Bottom) Fractional difference between the data and the fit. Observe the relation is universal to roughly 20\%. This means that the hexadecapole moment can be approximately expressed in terms of just the stellar mass, spin and quadrupole moment. \label{fig:M4-M2-intro} } \end{figure} \subsection{Executive Summary} Given the length of the paper, let us here present a brief summary of the main results. First, we confirm that the LORENE and RNS codes lead to numerically extracted multipole moments up to hexadecapole order that are not only consistent with each other, but also consistent with the slow-rotation approximation. We find that the numerical codes become inaccurate in the extraction of the mass hexadecapole moment for $\chi \lesssim 0.1$--$0.2$, while the slow-rotation expansion becomes inaccurate for $\chi \gtrsim 0.3$, where $\chi \equiv S_1/M_0^{2}$, with $S_1$ the spin angular momentum and $M_0$ is the stellar mass. The latter suggest that spin corrections to the moment of inertia and the quadrupole moment become important only for NSs with spin frequency larger than 100-450Hz, depending on the stellar compactness. Second, we confirm the results of~\cite{Pappas:2013naa,Chakrabarti:2013tca}, who found that the moment of inertia and the mass quadrupole moment obey approximately EoS-independent relations for arbitrarily rapidly spinning NSs, when one parametrized the stellar sequence in terms of $\chi$. We further confirm that the current octupole and the mass quadrupole moments also satisfy approximately EoS-independent relations to roughly 10\% variability~\cite{Pappas:2013naa}. We find that these relations hold not only for NSs, but also for QSs. Third, we investigate for the first time whether approximately EoS-independent relations exist in full GR between the mass hexadecapole and the mass quadrupole moment. The top panel of Fig.~\ref{fig:M4-M2-intro} shows the reduced (dimensionless) mass hexadecapole as a function of the mass quadrupole (see Eq.~\eqref{eq:def-dimensionless} for a definition), for various realistic NS and QS EoSs and spin frequencies. For comparison, the top panel also shows an analytic fit to all the data and the Newtonian relations of~\cite{Stein:2013ofa} for an $n=0.5$ polytropic EoS. The bottom panel shows the fractional difference between all the data and the fit. Observe that the relation is both approximately EoS-independent and also spin-independent, with 20\% variability at most. The variability though drops to only 10\% if the sequence of NSs and the sequence of QSs were considered seperately. Fourth, we find that the approximate EoS-universality between the hexadecapole and the quadrupole moments is worse than that found between the octupole and the quadrupole moments. This suggests that the universality becomes worse as one considers higher-$\ell$ multipoles, which is consistent with the Newtonian limit~\cite{Stein:2013ofa}. The top panel of Fig.~\ref{fig:M4-M2-intro} shows that the full GR $\bar{M}_{4}$--$\bar{M}_{2}$ relation approaches the Newtonian result rapidly as $\bar{M}_{2}$ increases, i.e.~as the compactness decreases. Fifth, we find that the Newtonian-limit relations of~\cite{Stein:2013ofa} are quite good at approximating the full GR relations, even for NSs with moderately large compactness. The vertical line in the top panel of Fig.~\ref{fig:M4-M2-intro} shows the value of $\bar{M}_{2}$ at which the Newtonian expression differs from the full GR one by 10\%. This occurs at $\bar{M}_2 = 3.5$, which corresponds to a compactness of approximately $0.2$. When considering the current octupole-mass quadrupole relation, the threshold value of $\bar{M}_2$ increases by an order of magnitude. This suggests that the Newtonian-limit of the universal relations become better approximations to the full GR relations as one considers relations between quadrupole and higher-$\ell$ multipole moments. Sixth, we confirm that the NS quadrupole and octupole moments scale with the spin-squared and the spin-cubed, respectively, and the coefficients roughly depend only on the mass and EoS~\cite{pappas-apostolatos}. We find that such a property is preserved also for QSs. We also find that the NS and QS hexadecapole moment scales with the fourth power of spin, and again, the coefficient only depends on the star's mass and EoS, extending the results found in~\cite{pappas-apostolatos}. Due to the universality mentioned in the paragraphs above, these coefficients are intimately related with each other. Such scaling is similar to that found for Kerr BHs, and may help as a guide to construct an exact analytic solution with a small number of parameters, by using one of the generating techniques to describe a NS and QS exterior spacetime. For example, the two-soliton solution of~\cite{1995JMP....36.3063M} was proposed as a possible candidate to describe the exterior spacetime of relativistic stars analytically~\cite{Pappas:2012nt}. However, such a solution has a hexadecapole moment whose spin dependence starts at quadratic order and not at quartic order. The scaling results suggest that a better analytic solutions would be one that respects the scaling of multipole moments described above. Seventh, we study the implications of our results to future X-ray observations of NSs, including NICER and LOFT. Reference~\cite{Baubock:2012bj} showed that the quadrupole moment affects the atomic line profiles significantly. Psaltis et al.~\cite{Psaltis:2013fha} found that the quadrupole moment and the stellar eccentricity contribute to the X-ray pulse profiles by 1--5\% and 10--30\% respectively, for pulsars with mass $\sim 1.8M_\odot$ and a spin frequency of 600Hz. Since the goal of NICER and LOFT is to measure the mass and radius independently within 5\% accuracy, both quantities must be included when analyzing pulse profiles. Reference~\cite{Psaltis:2013fha}, however, carried out this analysis in the slow-rotation approximation, neglecting cubic and higher order terms in the spin. Our construction of NS and QS solutions to quartic order in spin allows us to estimate the systematic errors in the analysis of~\cite{Psaltis:2013fha} due to the quadratic order in spin truncation. Quartic order in spin terms should lead to order $\chi^2$ correction to quadratic order in spin effects. The fastest spinning pulsar~\cite{2006Sci...311.1901H} discovered could have $\chi \sim 0.53$, which means that quartic in spin terms could lead to corrections of $\sim 30\%$ on quadratic in spin effects. Given that the leading-order-in-spin contribution of the quadrupole moment affects the X-ray pulse profiles by 1--5\%, the spin correction to the quadrupole moment can be neglected in future X-ray observations; one expects similar results to hold for the effect of the hexadecapole moment on the pulse profile. The quartic in spin corrections to the stellar eccentricity, however, can lead to modifications on the pulse profile of $\sim 6\%$ for rapidly rotating pulsars. Such spin corrections as a function of the stellar compactness are EoS-insensitive if one restricts the star to be a NS and not a QS, with stellar compactness larger than 0.15. We also found that although the (leading-order in spin) eccentricity-compactness relation is approximately EoS independent within $\sim 1\%$ for NSs, as found in~\cite{Baubock:2013gna} and used in~\cite{Psaltis:2013fha}, this independence breaks for QSs. For the latter, different QS EoSs lead to a variability of order 10\%. Such relatively large variability originates from the relation between the moment of inertia and the compactness. This means that if one wants to use the eccentricity-compactness or the moment-of-inertia-compactness relation to achieve 5\% accuracy in the measurement of the mass and radius, one needs to assume that the pulsar is a NS. The new universal relations found in this paper should allow us to reduce the number of parameters in X-ray observations, breaking degeneracies and improving parameter estimation. \subsection{Organization and Convention} The organization of the paper is as follows. In Sec.~\ref{sec:slow-rot}, we explain how to construct perturbative NS and QS solutions to quartic order in spin in the slow-rotation approximation, extending the results of~\cite{hartle1967,benhar}. We present the perturbed Einstein equations, the asymptotic behavior of the solution near the stellar center, the exterior solutions, and the matching conditions at the stellar surface and explain how to extract multipole moments. In Sec.~\ref{sec:rapid-rot}, we explain how to construct rapidly-rotating NS and QS solutions using the LORENE and RNS codes. In Sec.~\ref{sec:results}, we present our numerical results. We show the universal relations among the multipole moments and compare them with the Newtonian relations. In Sec.~\ref{sec:X}, we explain how our work may help to estimate the pulse profile of X-ray observations. We conclude in Sec.~\ref{sec:conclusions} and point to future work. Henceforth, we use geometric units with $c=1=G$. The $\mathcal{O}(N)$ symbol will represent a term of order $N$ or a term that is $\sim N$ in magnitude. Given the length of the paper, we find it convent to summarize the following conventions here. We use the following masses: \begin{itemize} \item $M_0$ is the Geroch-Hansen mass monopole moment~\cite{geroch, hansen}. \item $M_$ is the stellar mass for a non-rotating star, which agrees with the Geroch-Hansen mass monopole moment in the slow-rotation limit. \item $M_\mrm{Komar}$ is the Komar mass, which agrees with $M_0$. \end{itemize} We use the following radial quantities: \begin{itemize} \item $R_\mrm{eq}$ is the gauge-invariant circumferential stellar radius at the equator. \item $R_\mrm{pol}$ is the gauge-invariant circumferential stellar radius at the poles. \item $R_$ is the stellar radius for a non-rotating star. \item $r_\mrm{eq}$ is the coordinate equatorial radius in quasi-isotropic coordinates [Eq.~\eqref{eq:quasi-isotropic}] (not to be confused with that in the Hartle-Thorne coordinates [Eq.~\eqref{Eq:metric-slow-rot}]). \end{itemize} We use the following spin-related quantities: \begin{itemize} \item $f$ is the stellar (linear) spin frequency. \item $\Omega$ is the stellar (angular) spin velocity, which is related to the linear spin frequency by $\Omega = 2 \pi f$. \item $S_1$ is the Geroch-Hansen current dipole moment~\cite{geroch, hansen}. \item $J$ is the magnitude of the spin angular momentum to leading order in spin frequency. \item $J_\mrm{Komar}$ is the Komar angular momentum, which agrees with $S_1$. \item $j$ is the dimensionless spin parameter, defined by $j = J/M_*^2$ and valid to leading order in spin frequency. \item $\chi$ is the dimensionless spin parameter, defined by $\chi = S_1/M_0^2$ and valid to all orders in spin frequency. \end{itemize}	\label{sec:conclusions} We investigated the existence of approximate EoS-universal, no-hair like relations for NSs and QSs, i.e.~relations between high and low-order multipole moments that are approximately independent of the EoS. We calculated NS and QS solutions and their multipole moments to $\ell=4$ order in three different ways: using two fully relativistic numerical codes valid, in principle, for arbitrarily fast spinning stars, and using a slow-rotation approximation to quartic order in spin. Such a calculation allowed us to confirm that the LORENE and RNS codes are indeed consistent with each other, in terms of the calculation of multipole moments. We were also able to establish that these codes become highly inaccurate for computations of higher multiple moments as the spin decreases below $\chi = 0.1$--$0.2$ due to numerical noise. We found that the $\bar{M}_4$--$\bar{M}_2$ relation only depends very weakly on the EoSs and spins, just like the $\bar{S}_3$--$\bar{M}_2$ relation found in~\cite{Pappas:2013naa}. Therefore, the NS and QS current octupole and mass hexadecapole moments can be expressed in terms of the first three moments to $\mathcal{O}(10\%)$ accuracy. However, the universality is weaker for the former compared to the latter. This suggests that universal relations may not exist for higher order $\ell$ modes, or at the very least, that they will deteriorate with increasing $\ell$ number. This is consistent with leading-order Newtonian results in a weak-field expansion~\cite{Stein:2013ofa}. We also found that the $\bar{M}_4$--$\bar{M}_2$ relation is very close to these Newtonian results, even for highly relativistic NSs with large compactness. Our results have both theoretical and practical applications. On the theoretical side, it may not be possible to find mathematical theorems that prove the existence of truly EoS-independent no-hair like theorems for NSs or QSs. Of course, the relations found have always been of an approximate nature, but we have here found evidence that the universality seems to break as one considers higher $\ell$ multipoles. On a practical side, the relations found here may be important in the measurement of the mass and radius of NSs with future X-ray observations, including NICER and LOFT. In particular, the hexadecapole moment and the quadratic spin corrections to the quadrupole moment and the stellar eccentricity may dominate the error budget of such measurements over statistical error. The universal relations found here should allow us to reduce the number of parameters and break degeneracies among them. It would be important to investigate if one could come up with relations with less EoS- and spin-variation by changing the normalization of the multipole moments. For example, one could define new dimensionless multipole moments via $\bar{M}^\mrm{new}_\ell \equiv (-1)^{\ell/2} M_\ell/(M_0^{\ell +1} \chi^\ell C^k)$ and $\bar{S}^\mrm{new}_\ell \equiv (-1)^{(\ell-1)/2} S_\ell/(M_0^{\ell +1} \chi^\ell C^k)$, and find the $k$ that minimizes the variation. Since some of the calculations performed in this paper are order of magnitude estimates, detailed calculations would be desirable, where one solves the null geodesic equations for the NSs and QSs (valid to quartic order in spin) constructed here, using a ray-tracing algorithm~\cite{Baubock:2011ke} and evaluate the impact of multipole moments, stellar eccentricity and spin corrections on the X-ray pulse profile. Reference~\cite{I-Love-Q-Science,I-Love-Q-PRD} showed that the universality holds not only between the stellar moment of inertia and quadrupole moment, but also among the quadrupolar tidal Love number~\cite{love,hinderer-love,damour-nagar,binnington-poisson}. Reference~\cite{Yagi:2013sva} found that the universality holds among the higher-$\ell$ tidal Love numbers. These results, together with the ones shown in this paper, suggest that universal relations also exist among higher-$\ell$ rotation-induced multipole moments and higher-$\ell$ tidal Love numbers. A straightforward extension of this work would be to investigate the relations among multipole moments at even higher $\ell$ order, such as $\ell=5$ and 6. It would be interesting to see if the EoS-universality indeed deteriorates, as we found here for lower $\ell$ modes. It would also be interesting to see if the higher-$\ell$ relations approach the Newtonian limit faster with increasing $\ell$. In order to achieve this goal, one would have to extend the slow rotation expansion to higher order in spin, or alternatively, attempt to extract higher-order multipole moments with LORENE or RNS. The latter may be feasible for rapidly-rotating stars with spin frequencies near the mass-shedding limit, but it would be extremely difficult for slowly-spinning stars due to numerical noise. Perhaps, what would be even more interesting is to understand physically why the EoS-universality is realized for the low-order multipoles in the first place. One way to understand this would be to investigate which part of the EoS is most responsible for the universality. Another approach would be to consider the Newtonian limit again, where one can tackle the problem (semi-)analytically, and break some of the approximations used in~\cite{Stein:2013ofa}. In particular, it would be interesting to investigate how breaking the elliptical isodensity approximation of~\cite{Lai:1993ve} impacts the EoS universality. Work along these lines is currently in progress. An analytic model to describe the NS and QS exterior spacetimes may be useful in astrophysical observations. References~\cite{Stute:2003aj,berti-stergioulas} calculated a three-parameter solution, based on~\cite{Manko:2000ud}, by using the formalism developed by Ernst~\cite{Ernst:1967wx,Ernst:1967by}; they found that such a solution describes nicely the exterior spacetime of a rapidly-rotating NS. This solution cannot accurately capture the features of the exterior spacetime of a slowly-rotating NS because its quadrupole moment does not vanish in the non-rotating limit. Reference~\cite{Pappas:2012nv} extended these studies by considering a four-parameter solution (the two-soliton solution), found in~\cite{1995JMP....36.3063M}, which includes the three-parameter solution of~\cite{Manko:2000ud} as a special case. Since this solution depends on four free parameters, it can describe the stellar exterior spacetime to octupole order~\cite{Pappas:2012nt}. However, its hexadecapole moment in general has a term that depends on the spin-squared, which is absent in NS and QS spacetimes, as we proved in this paper. The slowly-rotating stellar solution valid to quartic order in spin found in this paper allows one to analytically express the stellar exterior spacetime with the correct hexadecapole moment. It would be interesting to extend these studies to find an analytic exterior spacetime for NSs and QSs valid to hexadecapole order and higher order without using the slow-rotation expansion. When calculating the multipole moments, we assumed that the stars are uniformly-rotating and unmagnetized. Newly-born NSs and hypermassive NSs formed after NS binary mergers are expected to be differentially rotating, and magnetars can have magnetic fields as large as $10^{17}$G. Therefore, it may be interesting to relax the uniformly-rotating and unmagnetized assumptions. To relax the former, one could use a given rotation law, such as that in~\cite{komatsu_eh1989}. To relax the latter, one could consider magnetic field configurations similar to those in~\cite{I-Love-Q-B}. Probably, the magnetic field will become progressively more important, as one considers higher-$\ell$ multipole moments. This is because the rotationally-induced multipole moments should become smaller in magnitude as one increases $\ell$. One could also study how differential rotation and magnetic field affect the relations among multipole moments in the Newtonian limit, namely extending the work in~\cite{Stein:2013ofa}.	14	3	1403.6243	Astrophysical charge-free black holes are known to satisfy no-hair relations through which all multipole moments can be specified in terms of just their mass and spin angular momentum. We here investigate the possible existence of no-hair-like relations among multipole moments for neutron stars and quark stars that are independent of their equation of state. We calculate the multipole moments of these stars up to hexadecapole order by constructing uniformly rotating and unmagnetized stellar solutions to the Einstein equations. For slowly rotating stars, we construct stellar solutions to quartic order in spin in a slow-rotation expansion, while for rapidly rotating stars, we solve the Einstein equations numerically with the LORENE and RNS codes. We find that the multipole moments extracted from these numerical solutions are consistent with each other and agree with the quartic-order slow-rotation approximation for spin frequencies below roughly 500 Hz. We also confirm that the current dipole is related to the mass quadrupole in an approximately equation-of-state-independent fashion, which does not break for rapidly rotating neutron stars or quark stars. We further find that the current-octupole and the mass-hexadecapole moments are related to the mass quadrupole in an approximately equation-of-state-independent way to roughly O(10%), worsening in the hexadecapole case. All of our findings are in good agreement with previous work that considered stellar solutions to leading order in a weak-field, Newtonian expansion. In fact, the hexadecapole-quadrupole relation agrees with the Newtonian one quite well even in moderately relativistic regimes. The quartic in spin, slowly rotating solutions found here allows us to estimate the systematic errors in the measurement of the neutron star's mass and radius with future x-ray observations, such as Neutron star Interior Composition ExploreR (NICER) and Large Observatory for X-ray Timing (LOFT). We find that the effect of these quartic-in-spin terms on the quadrupole and hexadecapole moments and stellar eccentricity may dominate the error budget for very rapidly rotating neutron stars. The new universal relations found here should help to reduce such systematic errors.	false	[ "neutron stars", "quark stars", "Neutron star Interior Composition ExploreR", "X-ray Timing", "stellar solutions", "future x-ray observations", "hexadecapole order", "such systematic errors", "quartic order", "RNS codes", "multipole moments", "solutions", "leading order", "stellar eccentricity", "hexadecapole", "rapidly rotating stars", "slowly rotating stars", "Large Observatory", "angular momentum", "moments" ]	5.464849	2.223879	-1
12311913	[ "Miller, Brendan P.", "Gallo, Elena", "Greene, Jenny E.", "Kelly, Brandon C.", "Treu, Tommaso", "Woo, Jong-Hak", "Baldassare, Vivienne" ]	2015ApJ...799...98M	[ "X-Ray Constraints on the Local Supermassive Black Hole Occupation Fraction" ]	127	[ "Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA ; Department of Physics and Astronomy, Macalester College, Saint Paul, MN 55105, USA", "Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA", "Department of Astrophysics, Princeton University, Princeton, NJ 08544, USA", "Physics Department, University of California, Santa Barbara, CA 93106, USA", "Physics Department, University of California, Santa Barbara, CA 93106, USA", "Astronomy Program, Department of Physics and Astronomy, Seoul National University, Seoul, Korea", "Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA" ]	[ "2014ApJ...787L..30R", "2014ApJ...791..133B", "2014ApJ...794....9M", "2014ApJ...794...34N", "2015ApJ...809L..14B", "2015ApJ...811...26T", "2015MNRAS.447.2772R", "2015MNRAS.451.3615N", "2015MNRAS.453..775G", "2016ApJ...817...20M", "2016ApJ...819..162P", "2016ApJ...823..112P", "2016ApJ...825..139P", "2016ApJ...827...58S", "2016ApJ...831..203P", "2016MNRAS.455..859S", "2016MNRAS.456.2537R", "2016MNRAS.460.2979V", "2016MNRAS.463..348V", "2016MNRAS.463..811O", "2016PASA...33...32V", "2016PASA...33...54R", "2017ApJ...835..223S", "2017ApJ...837...48C", "2017ApJ...839..120W", "2017ApJ...841...51F", "2017ApJ...842...72Y", "2017ApJ...848...92P", "2017ApJ...850...15P", "2017ApJ...851..106F", "2017IJMPD..2630021M", "2017MNRAS.468.3935H", "2017MNRAS.471.4286F", "2017MNRAS.472.1526C", "2017NewAR..79...59X", "2017PASJ...69...90T", "2017PhRvD..96j3003A", "2018ASPC..517..719P", "2018ApJ...858...20V", "2018ApJ...858..118N", "2018ApJ...861...39D", "2018ApJ...861..142C", "2018ApJ...864L...6P", "2018ApJ...868..152B", "2018MNRAS.473.5698D", "2018MNRAS.478.2576M", "2018MNRAS.481.3278R", "2018arXiv181006814P", "2019Ap&SS.364..165T", "2019ApJ...872..104N", "2019ApJ...874...77L", "2019ApJ...874..117B", "2019ApJ...878...18S", "2019ApJ...883L..18G", "2019ApJ...884...78L", "2019ApJ...885L...3H", "2019BAAS...51c.175B", "2019MNRAS.482.2913B", "2019MNRAS.484..794G", "2019MNRAS.488..685M", "2019MNRAS.489.1006R", "2019MNRAS.489.5413M", "2019arXiv190308670G", "2020A&ARv..28....4N", "2020ARA&A..58..257G", "2020ApJ...888...36R", "2020ApJ...892...93C", "2020ApJ...896...10B", "2020ApJ...897..103S", "2020ApJ...898..106H", "2020ApJ...898..164K", "2020MNRAS.492.2528L", "2020MNRAS.493..899H", "2020MNRAS.493L..76K", "2020MNRAS.495..123Q", "2020MNRAS.496.4061D", "2020MNRAS.498.2219V", "2020SSRv..216...39M", "2021A&A...651A.109D", "2021ApJ...911..134K", "2021ApJ...914...69C", "2021ApJ...917...17G", "2021ApJ...922L..40L", "2021MNRAS.503..679M", "2021MNRAS.503.3568K", "2021MNRAS.503.6098R", "2021MNRAS.505.5129B", "2021NatRP...3..732V", "2022ApJ...929..147D", "2022ApJ...936...82S", "2022ApJ...940..111U", "2022MNRAS.509.2920N", "2022MNRAS.511.2631A", "2022MNRAS.511.4109D", "2022MNRAS.514.4912H", "2022MNRAS.515.4038T", "2022MNRAS.516.2112K", "2022NatAs...6...26R", "2022NatAs...6.1452A", "2023A&A...676A..38E", "2023A&A...677A.123I", "2023ApJ...946...51C", "2023ApJ...953..142C", "2023ApJ...954L...4K", "2023ApJ...955...78C", "2023ApJ...955L...6Y", "2023ApJ...958..115K", "2023MNRAS.518.1880B", "2023MNRAS.518.4672S", "2023MNRAS.519.5848T", "2023MNRAS.520.3916D", "2023MNRAS.521.2845C", "2023MNRAS.525.1182S", "2023arXiv230600898S", "2023arXiv230715111C", "2024A&A...681A..97A", "2024A&A...682A..36H", "2024AJ....167...31O", "2024ApJ...960...15K", "2024ApJ...964...11E", "2024MNRAS.527.4643M", "2024OJAp....7E..26T", "2024arXiv240319635B", "2024arXiv240408910F", "2024arXiv240509441C", "2024arXiv240601707S", "2024arXiv240605778D" ]	[ "astronomy" ]	18	[ "black hole physics", "galaxies: nuclei", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1982MNRAS.200..115S", "1997AJ....114.2366C", "1997MNRAS.286L..50D", "1998ApJ...492..554N", "2000ApJ...539L..13G", "2001A&A...372...29G", "2001AJ....122.2469G", "2001Sci...293.1116M", "2002MNRAS.335..965Y", "2004ApJ...604L..89H", "2005ApJ...628..137V", "2006ApJ...640..143S", "2006MNRAS.365...11C", "2006MNRAS.372...21A", "2007ApJ...655..144M", "2007ApJ...665.1489K", "2007ApJ...667..117T", "2008A&A...486..119B", "2008ApJ...680..154G", "2009ApJ...698..198G", "2009ApJ...699..626H", "2009ApJ...699..800V", "2009MNRAS.400.1911V", "2010A&A...518L.155F", "2010A&ARv..18..279V", "2010AJ....140..546W", "2010ApJ...714...25G", "2010ApJ...714..713S", "2010ApJ...716..269W", "2010ApJ...717..640P", "2010ApJ...720..368X", "2010ApJS..189..270C", "2010MNRAS.402..673B", "2010MNRAS.408.1139V", "2010MNRAS.409..500O", "2011ApJ...726...34C", "2011ApJ...730..145V", "2011ApJ...740L..42D", "2011ApJ...742...68J", "2011ApJS..195...10X", "2011Natur.470...66R", "2011Natur.474..616M", "2012A&A...537L...8C", "2012A&A...545L...6S", "2012AdAst2012E..15N", "2012ApJ...745L..13M", "2012ApJ...746...90A", "2012ApJ...747...57M", "2012ApJ...748..124Y", "2012ApJ...752...46L", "2012ApJ...756L..19W", "2012MNRAS.419..267P", "2012MNRAS.419.2497B", "2012MNRAS.422.1690P", "2012MNRAS.422.2051N", "2012NatCo...3.1304G", "2012RPPh...75l4901V", "2012Sci...337..544V", "2013A&A...558A..14M", "2013ApJ...764..184M", "2013ApJ...766..121M", "2013ApJ...767...13S", "2013ApJ...771..116J", "2013ApJ...772...49W", "2013ApJ...773..150S", "2013ApJ...775..116R", "2013ApJ...777..133M", "2013ApJ...778...47S", "2013ApJ...778..130T", "2013MNRAS.428..421S", "2013MNRAS.432..530R", "2013MNRAS.433.1607L", "2013MNRAS.434.2600F", "2013MNRAS.436.2576L", "2013MNRAS.436.2989L", "2014ARA&A..52..529Y", "2014ApJ...780....9W", "2014ApJ...782...55Y", "2014ApJ...782...69L", "2014ApJ...783L..33K", "2014ApJ...784L..38M", "2014ApJ...787L..30R", "2014ApJ...791..133B", "2014ApJ...794....9M", "2014ApJ...794...34N", "2014GReGr..46.1702N", "2014MNRAS.441.1059V", "2014MNRAS.442.2036D", "2014MNRAS.442.2751T", "2014MNRAS.442.3616L", "2014MNRAS.443..648A", "2014MNRAS.443.2410F", "2014Natur.513..398S" ]	[ "10.1088/0004-637X/799/1/98", "10.48550/arXiv.1403.4246" ]	1403	1403.4246_arXiv.txt	Observations of high-redshift quasars indicate that supermassive black holes (SMBHs\footnote{We use the term ``supermassive'' to indicate masses of $M_{\rm BH}\simgt3\times10^{5} M_{\odot}$ for central black holes, as in Greene (2012).}) are already present in the early universe (e.g., Vestergaard \& Osmer 2009; Willott et al.~2010; Mortlock et al.~2011). SMBHs with $M_{\rm BH}\simgt10^{9} M_{\odot}$ by $z\simgt6$ are extremely challenging to grow from Population III remnants (``light'' seeds of $\sim100 M_{\odot}$; e.g., Whalen \& Fryer 2012; Madau et al.~2014; Taylor \& Kobayashi 2014), but can derive from direct gas collapse precursors (``heavy'' seeds of $\sim10^{5} M_{\odot}$; e.g., Begelman 2010; Johnson et al.~2013; Ferrara et al.~2014). However, the unresolved cosmic \hbox{X-ray} background implies SMBHs are not common (or else are generally quasi-quiescent) in high-redshift galaxies (Salvaterra et al.~2012), a possibility also suggested by stringent constraints on average nuclear X-ray luminosities obtained from stacking deep field {\it Chandra\/} observations (Triester et al.~2013). For typical expected subsequent black hole growth (Shankar et al.~2013), and in line with the SMBH mass function derived from broad-line quasars (Natarajan \& Volonteri 2012), these results may be more consistent with sparse heavy seeding than with slow initial growth of omnipresent light seeds. Despite significant and ongoing theoretical and observational advances, the particular seed mechanism predominantly responsible for SMBH formation is not yet conclusively established (see reviews by Volonteri 2012; Volonteri \& Bellovary 2012; Natarajan 2014; and references therein). The evolution of SMBHs appears to be entwined with that of their host galaxies. This is suggested by the $M_{\rm BH}-\sigma$ relation linking the central black hole mass to the bulge stellar velocity dispersion, which holds for both quiescent spheroids (G{\"u}ltekin et al.~2009; McConnell \& Ma~2013) and active galactic nuclei (Woo et al.~2010, 2013) and may be redshift-dependent (Treu et al.~2007; Lapi et al.~2014). SMBH feedback provides one plausible linking mechanism (Sun et al.~2013), as predicted by semi-empirical modeling (Croton et al.~2006; Shankar et al.~2013) and in a few cases now directly measured (e.g., Feruglio et al.~2010; Cano-D{\'{\i}}az et al.~2012; Liu et al.~2013). Mergers and intermittently efficient accretion in larger SMBHs spur growth and remove observational signatures of their birth, but smaller SMBHs have more subdued histories and undergo mostly secular evolution (Jiang et al.~2011). Consequently, both the mass distribution and the very rate of occurance of SMBHs in lower-mass galaxies contain archaeological information on the initial seed formation mechanism. A robust conclusion from semi-analytical modeling is that smaller galaxies are more likely to contain SMBHs when Pop~III remnants, rather than direct gas collapse, provide the dominant\footnote{Intermediate mass seeds, for example from nuclear star cluster collapse (Davies et al.~2011; Lupi et al.~2014), are a third possibility.} seeding mode (Volonteri \& Natarajan 2009; Volonteri 2010; van Wassenhove et al.~2010). This is because cold low-metallicity gas is only able to collapse to a central massive object in halos with low spin parameter, otherwise disk fragmentation leads to star formation (van Wassenhove et al.~2010). The fraction of halos forming such heavy seeds should exceed 0.001 to produce SMBHs at $z=6-7$ (Petri et al.~2012). Using a First Billion Years cosmological hydrodynamical simulation, Agarwal et al.~(2014) identify several pristine\footnote{Enriched gas cannot directly collapse to produce a massive seed (e.g., Ferrara et al.~2013).} atomic-cooling haloes that could host direct-collapse massive seeds, and note that these haloes are universally close to protogalaxies and exposed to a high flux of Lyman-Werner radiation (as also found by, e.g., Latif et al.~2013a, b; Dijkstra et al.~2014). Measurement of the occupation fraction (i.e., the percentage of galaxies hosting SMBHs) in nearby galaxies, particularly at low stellar masses $M_{\rm star}<10^{9-10} M_{\odot}$, is an effective {\it observational\/} discriminator between light versus heavy seeds (Greene 2012). The limited $\simlt$10$^{8}$ yr lifetime of luminous quasars suggests (Soltan 1982; Yu \& Tremaine 2002), consistent with observations, that the most massive ``inactive'' galaxies invariably host SMBHs now accreting/radiating only at $\simlt10^{-5}$ Eddington, but the occupation fraction in lower mass galaxies remains uncertain. Clearly some low-mass galaxies do possess SMBHs\footnote{The masses of central black holes in dwarf galaxies are difficult to measure precisely, but the following examples are likely near or above our adopted definitional threshold for a SMBH.}, even active ones. For example, the dwarf galaxy Henize 2-10 hosts an accreting SMBH as revealed by \hbox{X-ray} and radio emission (Reines et al.~2011), and features central blue clumps of star-formation within a red early-type system (Nguyen et al.~2014). Mrk 709 is an interacting pair of dwarfs, the Southern of which has a central \hbox{X-ray} and radio source indicating the presence of a SMBH (Reines et al.~2014). Within the {\it Chandra\/} Deep Field South Survey, Schramm et al.~(2013) identify three galaxies with $M_{\star}<3\times10^{9} M_{\odot}$ that have \hbox{X-ray} emitting SMBHs. Yuan et al.~(2014) describe four dwarf Seyferts with $M_{\rm BH}\simlt10^{6} M_{\odot}$, two of which are detected in X-rays with $L_{\rm X}\sim10^{41}$~erg~s$^{-1}$. A sample of 151 dwarf galaxies with candidate SMBHs as identified from optical emission line ratios and/or broad H$\alpha$ emission is presented by Reines et al.~(2013; see also references therein). The ultra-compact dwarf galaxy M60-UCD1 is indicated by a central velocity dispersion peak to have a SMBH with $M_{\rm BH}=2.1\times10^{7} M_{\odot}$, but here the large black hole mass fraction suggests substantial stellar mass has been stripped from the galaxy (Seth et al.~2014). For each example of a low-mass galaxy that has observational evidence for a central SMBH, there are 10--100 similar galaxies for which the presence or absence of a black hole is currently impossible to measure. However, dynamical mass constraints are quite stringent for some Local Group objects (the spiral M33: Gebhardt et al.~2001; Merritt et al.~2001; the dwarf elliptical NGC~205: Valluri et al.~2005), which effectively rules out a 100\% SMBH occupation fraction. High spatial resolution \hbox{X-ray} observations can efficiently identify very low-level SMBH activity (Soria et al.~2006; Pellegrini 2010) without contamination from the stellar emission that dilutes optical searches. Nuclear \hbox{X-ray} emission directly measures high-energy accretion-linked radiative output and additionally serves as a plausible proxy for mechanical feedback (Allen et al.~2006; Balmaverde et al.~2008). \hbox{X-ray} studies of low-level SMBH activity are best conducted on galaxies with low star formation rates to eliminate potential contamination from high-mass \hbox{X-ray} binaries. For statistical purposes the sample must span a wide range in $M_{\rm star}$ and be unbiased with respect to optical or \hbox{X-ray} nuclear properties. These criteria are satisfied by the AMUSE\footnote{AGN Multiwavelength Survey of Early-Type Galaxies}-Virgo (Gallo et al.~2008, 2010; G08, G10 hearafter) and AMUSE-Field (Miller et al.~2012a, 2012b; M12a, M12b hereafter) surveys, which are Large {\it Chandra\/} Programs that together targeted 203 optically-selected early-type galaxies at $d<30$~Mpc, and now include {\it HST\/}, {\it Spitzer\/}, and {\it VLA\//JVLA} coverage. Almost all of these galaxies have $L_{\rm X}<10^{41}$~erg~s$^{-1}$ and $L_{\rm X}/L_{\rm Edd}<10^{-5}$, below limits commonly used to distinguish active galactic nuclei (AGN) from ``inactive'' galaxies. In this work we use the AMUSE dataset to obtain the first simultaneous constraints upon the scaling of nuclear activity with host galaxy stellar mass and the local supermassive black hole occupation fraction, and derive guidelines for the precision that may be achieved with a larger sample or a next-generation \hbox{X-ray} telescope.	Our simultaneous fitting of the SMBH occupation fraction and the scaling of nuclear \hbox{X-ray} luminosity with stellar mass constrains SMBHs to be present in $>$20\% of early-type galaxies with $M_{\rm star}<10^{10} M_{\odot}$ and suggests the dependence of $\log{L_{\rm X}}$ upon $\log{M_{\rm star}}$ has a slope of $\sim$0.7--0.8. This work provides promising if inconclusive information on the local SMBH occupation fraction and also supports a downsizing trend in low-level SMBH activity. The highly sub-Eddington objects that make up the AMUSE dataset are expected to feature radiatively inefficient accretion flows (RIAFs). Bondi accretion of even the limited gas provided by stellar winds (Volonteri et al.~2011) near the nuclei of early-type galaxies would predict greater \hbox{X-ray} luminosities than observed; the efficiency as well as the accretion rate must be low in these objects (Soria et al.~2006; Ho 2009). Either an advection-dominated accretion flow (e.g., Di Matteo \& Fabian 1997; Narayan et al.~1998) or an outflow/jet component (e.g., Soria et al.~2006; Plotkin et al.~2012) is required. In general, the efficiency in these hot flows is theoretically expected to decrease toward lower accretion rates (Yuan \& Narayan 2014 and references therein). Although the Bondi radius is directly resolved by Chandra in deep observations of NGC~3115, the temperature profile is inconsistent with simple RIAF models (Wong et al.~2014). Fueling of a RIAF by steady-state stellar winds may be supplemented by intermittent processes such as tidal disruption, or by gradual stripping of central stars (e.g., MacLeod et al.~2013). While we cannot constrain the physical mechanism responsible for the observed \hbox{X-ray} emission, the simplest explanation for downsizing in low-level SMBH activity would be that the relative rate of accretion is higher in smaller galaxies, with a fixed low efficiency. We reiterate that the downsizing we identify here is restricted to low-level SMBH activity and may not apply to AGNs. The methodology we use here is flexible and could also be applied to deep surveys of AGN. For example, the 4~Ms CDFS contains active galactic nuclei including to relatively modest $M_{\rm star}\simlt3\times10^{9} M_{\odot}$ (Schramm \& Silverman 2013; Schramm et al.~2013) and at low levels of activity (Young et al.~2012) as well as normal galaxies at cosmological distances (Lehmer et al.~2012), and opens substantial additional volume albeit at lower sensitivity. We provide an illustration of applying this general technique to simulated deep field galaxies in Figure~5, and are pursuing this approach in Greene et al.~(in preparation). In this higher $L_{\rm X}$ regime we populate X-ray luminosities drawing from a uniform Eddington ratio distribution with a power-law slope of $-0.65$ as in Aird et al.~(2012), assuming $\log{M_{\rm BH}} = \log{M_{\rm star}}-2.8$. For a hypothetical combined CDFS+CDFN+AEGISXD sample with a typical detection sensitivity of $\log{L_{\rm X}}\simeq40$ out to $z<0.4$, we expect about 15000 $z<0.4$ galaxies (estimated from Cardamone et al.~2010; Xue et al.~2010) of which $\sim$300 or 2\% should host \hbox{X-ray} AGNs (estimated from Xue et al.~2010, 2011; Lehmer et al.~2012). We confirmed with an artificially large sample that the distribution of \hbox{X-ray} detections (color gradients in Figure~5) can be used to infer the occupation fraction. For example, the percentages of \hbox{X-ray} AGNs in hosts with $\log{M_{\rm star}}<9.5$ is 11.9\% with full occupation versus 6.1\% with half occupation. This is statistically distinguishable at 99\% confidence for the expected $\simeq$300 AGNs (black crosses in Figure~5) if the other model parameters are known or fixed by theory; full versus half occupation predicts 37 versus 17 \hbox{X-ray} AGNs in hosts with $\log{M_{\rm star}}<9.5$). This test will increase in power with the upcoming deeper CDFS exposure. We are also refining our measurement of occupation fraction within the AMUSE sample through incorporating the influence of large-scale environment upon low-level SMBH activity. The scaled nuclear \hbox{X-ray} luminosities of early-type galaxies apparently decrease from isolated to group to cluster environments (M12a,b). This may reflect greater quantities of cold gas in field galaxies (e.g., Oosterloo et al.~2010), for example due to reduced stripping relative to their cluster counterparts. Cold accretion has been inferred to be relevant to low-level SMBH activity from studies of brightest cluster galaxies (Russell et al.~2013) and dust (Martini et al.~2013), and AGNs preferentially inhabit gas-rich galaxies (Vito et al.~2014). The recent tentative finding that nuclear star clusters in massive early type galaxies are bluer in the field (B14) implies that field nuclear star clusters formed at lower metallicities and/or experienced more recent star formation, relative to cluster counterparts; this in turn suggests that cold gas can eventually filter down to the central regions where it is available (either directly or via enhanced star formation and stellar winds) to be heated and inefficiently accreted onto the central SMBH. We are continuing to investigate the impact of Mpc-scale densities using new {\it Chandra\/} observations of early-type galaxies located within cosmic voids. However, the analysis presented here is not biased because the slopes of the $L_{\rm X}-M_{\rm star}$ relation are consistent between the AMUSE-Field and AMUSE-Virgo samples (M12b); instead, the uncertainties are potentially slightly inflated. Including any environmental dependence, once quantified at high significance, will helpfully decrease the scatter in the $L_{\rm X}(M_{\rm star})$ relation in the combined AMUSE dataset. Additional multiwavelength information will provide better understanding of both individual objects and of the overall population (e.g., the distribution of galaxies showing radio, or optical, indications of nuclear activity; Reines et al.~2013). New dynamical mass measurements with a 30m class telescope would help clarify the mass distribution of SMBHs in smaller galaxies, providing a complementary probe of black hole birth and growth (van Wassenhove et al.~2010). In this context it is interesting that no galaxies with $M_{\rm star}<10^{10} M_{\odot}$ (without stripping; Seth et al.~2014) are yet known with confirmed $M_{\rm BH}>10^{6} M_{\odot}$. Tidal disruption transients, particularly from white dwarfs, can provide complementary insight into lower-mass SMBHs (Clausen \& Eracleous 2011; MacLeod et al.~2014). Pairing observational advances with increasingly sophisticated theoretical models will help discriminate between models of seed formation. \begin{figure} \includegraphics[scale=0.68]{deep_sim.ps} \figcaption{\small Distribution of AGN X-ray detections for mock deep field catalogs with 50\% and 100\% occupation fractions for $M_{\rm star}<10^{10} M_{\odot}$. The colors indicate detection density with the black crosses a realization with 15000 total galaxies and $\sim$300 X-ray AGN, $\pm$20 depending on the occupation fraction. The top histogram shows X-ray detected AGN for half (black) and full (gray) occupation.} \end{figure}	14	3	1403.4246	Distinct seed formation mechanisms are imprinted upon the fraction of dwarf galaxies currently containing a central supermassive black hole. Seeding by Population III remnants is expected to produce a higher occupation fraction than is generated with direct gas collapse precursors. Chandra observations of nearby early-type galaxies can directly detect even low-level supermassive black hole activity, and the active fraction immediately provides a firm lower limit to the occupation fraction. Here, we use the volume-limited AMUSE surveys of ~200 optically selected early-type galaxies to characterize simultaneously, for the first time, the occupation fraction and the scaling of L <SUB>X</SUB> with M <SUB>star</SUB>, accounting for intrinsic scatter, measurement uncertainties, and X-ray limits. For early-type galaxies with M <SUB>star</SUB> < 10<SUP>10</SUP> M <SUB>⊙</SUB>, we obtain a lower limit to the occupation fraction of >20% (at 95% confidence), but full occupation cannot be excluded. The preferred dependence of log L <SUB>X</SUB> upon log M <SUB>star</SUB> has a slope of ~0.7-0.8, consistent with the "downsizing" trend previously identified from the AMUSE data set, and a uniform Eddington efficiency is disfavored at ~2σ. We provide guidelines for the future precision with which these parameters may be refined with larger or more sensitive samples.	false	[ "full occupation", "X-ray limits", "direct gas collapse precursors", "SUB", "a higher occupation fraction", "the occupation fraction", "%", "dwarf galaxies", "~2σ", "the active fraction", "a firm lower limit", "measurement uncertainties", "a central supermassive black hole", "L", "intrinsic scatter", "even low-level supermassive black hole activity", "a lower limit", "the fraction", "Eddington", "log M <SUB>star</SUB" ]	14.701868	7.077017	119
437528	[ "Steinacker, J.", "Ormel, C. W.", "Andersen, M.", "Bacmann, A." ]	2014A&A...564A..96S	[ "Coreshine in L1506C - Evidence for a primitive big-grain component or indication for a turbulent core history?" ]	17	[ "UJF-Grenoble 1/CNRS-INSU, Institut de Planétologie et d'Astrophysique de Grenoble (IPAG) UMR 5274, 38041, Grenoble, France ; Max-Planck-Institut für Astronomie, Königstuhl 17, 69117, Heidelberg, Germany", "Astronomy Department, University of California, Berkeley, CA, 94720, USA", "UJF-Grenoble 1/CNRS-INSU, Institut de Planétologie et d'Astrophysique de Grenoble (IPAG) UMR 5274, 38041, Grenoble, France", "UJF-Grenoble 1/CNRS-INSU, Institut de Planétologie et d'Astrophysique de Grenoble (IPAG) UMR 5274, 38041, Grenoble, France" ]	[ "2014A&A...572A..20L", "2015A&A...582A..70S", "2015ApJ...806..196K", "2015ApJ...811...38W", "2015P&SS..116...64J", "2016A&A...588A..44Y", "2016PCCP...18.5159R", "2016arXiv160508032D", "2017A&A...605A..99T", "2018A&A...614A..95S", "2018PASJ...70...76O", "2021A&A...647A.109S", "2021ApJ...910...26T", "2023A&A...671A.111J", "2023A&A...676A...4C", "2024A&A...682A..61C", "2024AN....34530126Z" ]	[ "astronomy" ]	102	[ "dust", "extinction", "ISM: clouds", "infrared: ISM", "ISM: individual objects: L1506", "scattering", "Astrophysics - Astrophysics of Galaxies" ]	[ "1977ApJ...217..425M", "1984ApJ...285...89D", "1993A&A...280..617O", "1994ApJ...430..713W", "1996A&AS..117..393B", "2003ApJ...585L.131V", "2003ApJ...598.1017D", "2004RvMP...76..125M", "2005MNRAS.356..654G", "2007ARA&A..45..339B", "2008ApJ...683..238K", "2009A&A...502..845O", "2010A&A...511A...9S", "2010A&A...512A...3P", "2010ApJS..186..259R", "2010Sci...329.1622P", "2011A&A...532A..43O", "2011ApJ...742....7A", "2012A&A...541A.154P", "2012ApJ...760..147Q", "2013A&A...551A..38P", "2013A&A...559A.133Y", "2013ARA&A..51...63S" ]	[ "10.1051/0004-6361/201322117", "10.48550/arXiv.1403.3650" ]	1403	1403.3650_arXiv.txt	\label{intro} Star formation research has identified the densest parts of molecular clouds as the sites where stars form. But the entire process from forming the molecular cloud, through the formation of the cloud cores, its evolution while having no central object, to the final collapse is not understood \citep{2007ARA&A..45..339B}. Physically, we are facing a coupled problem of gas dynamics including turbulent motions, gravitation, and magnetic fields where cooling and heating of the gas and dust are influenced by the chemical processes, grain processing, and the external illumination. The theory of gravo-turbulent star formation predicts that molecular cloud cores form at the stagnation points of the complex turbulent flow pattern \citep{2004RvMP...76..125M}. But since supersonic turbulence does not create hydrostatic equilibrium configurations, the density structures are transient and dynamically evolving, as the different contributions to virial equilibrium do not balance \citep{2003ApJ...585L.131V}. Therefore, a statistical approach in collecting core properties has been a main approach to study their evolution. For individual cores, their history has been accessed by studying their chemical state \citep{2013A&A...551A..38P}. The filamentary Taurus star formation region \citep[see, e.g.,][]{2012arXiv1206.2115Q} is an ideal site to test these theories and to determine the properties of cores due to its proximity and coverage of all low-mass star formation stages. In a recent analysis of a Taurus filament fragment L1506C based on dust emission, and line emission of C$^{18}$O, N$_2$H$^+$, $^{13}$CO, and C$^{17}$O, \citet[][Pea10 hereafter]{2010A&A...512A...3P} show that the fragment combines interesting and surprising properties. Interpreted as a core-envelope system, the core located within 3$\times$10$^4$ au shows low densities n(H$_2)<5\times$10$^{10}$ m$^{-3}$, but high depletion $>$30 of C$^{18}$O and low temperatures 8-10 K depending on the tracer. Kinematic information was derived from line emission measured along a cut through the filament. The line-of-sight integrated emission was modeled with a 1D line transfer code. The authors argued that the line profiles show evidence for contraction of the core plus rotation with an infall speed of about 100 m/s. In turn, the velocity gradient of the envelope visible in the $^{13}$CO line was found to be opposite in direction to that of the core as traced by the C$^{18}$O line. Moreover, the core showed extremely low turbulence v$_{turb}$(FWHM)$<$68 m/s. Based on the contraction of the core and its low density, the authors concluded that the core is about to become a pre-stellar core, but has not yet reached the density criterion suggested by \citet{2008ApJ...683..238K} for a pre-stellar core. An important piece of evidence was added to this picture by \citet{2010Sci...329.1622P}. They reported on the detection of coreshine for about half of 110 investigated cores, among them L1506C. According to the analysis presented in \citet{2010A&A...511A...9S} for the core L183, the interstellar radiation can be scattered at wavelengths around 4 $\mu$m when the core gas hosts dust grains which are about a factor of 4 larger than the maximal silicate grain size of about 0.25 $\mu$m advocated by the MRN model. Investigating Herschel observations of the entire filament harboring L1506C, \citet{2013A&A...559A.133Y} found that the grain opacity has to increase across the filament to fit simultaneously the near-IR extinction and FIR Herschel emission profiles. They interpreted this change to be a consequence of the coagulation of dust grains to form fluffy aggregates with the grain average size being increased by a factor of 5. In this paper, we analyze the coreshine seen towards L1506C and discuss its implication for the underlying grain size distribution and the history of the core. In Sect.~2, we derive the range of coreshine surface brightness being consistent with the IRAC 3.6 $\mu$m data of L1506C. For a model core, we calculate the scattered light intensity for MRN-type size distributions as well as for time-dependent distributions arising from coagulation calculations in Sect.~3. Comparing the derived surface brightnesses we propose possible scenarios to account for the coreshine in Sect.~4.	\label{conclusions} We have investigated the scattered light observed from the central part of L1506C in the Spitzer channel 1 and put forward two scenarios to account for the coreshine surface brightness. Mixing the core case with a constant MRN dust size distribution using a size limit of 0.25 $\mu$m does not lead to sufficient scattered light to explain the observed surface brightness of cores with coreshine. The core may host a component of large grains with sizes up to 0.65 $\mu$m that has not been formed in the core rather than being a primitive component from earlier star formation cycles before the Taurus filaments came to being. Using a simple density structure, an ISRF based on the DIRBE map plus a local component from the star formation region, and standard dust opacities we were able to reproduce the coreshine level with our radiative transfer calculations. An implication of the existence of a global primitive large-grain component would be that most low-mass cores comparable to L1506C should show coreshine as far as the background level allows its detection, shielding effects do not block the radiation either during illumination or on its way to the observer, or large-scale processes like supernovae have modified the size distribution \citep{2011ApJ...742....7A,2012A&A...541A.154P}. A future publication will investigate this prospective. Another possibility is that the large grains responsible for the scattering have been formed in the core. To enable efficient grain growth, we maintained the turbulence on the largest scale for 1 Myr which is a basic assumption of the gravo-turbulent scenario for star formation. But even with the most optimistic grain growth model, and increasing either the turbulence or the density beyond the currently observed values, we failed to reproduce the observed coreshine within the limits of our simple model. However, our results indicate that grain growth in a core that is denser and more turbulent could lead to the observed coreshine surface brightness. Within the gravo-turbulent scenario, the same large-scale motions that would have created the density maximum in the filament at the location of L1506C could have torn the core partially apart leaving just the gravitationally bound part. The question is if the velocity gradients that have been inferred from the emission line analysis along a cut in Pea10 are in agreement with this scenario. Since turbulence on the largest scales decays slowest in the Kolmogorov picture, it is on the large-scales that remnants of such a tearing motion should still be seen. The observed outward motion of the envelope at large scales can be interpreted in this way. It is more complex to access the present kinematical signatures on the smaller scales, that is on the scale of the core. First, the turbulence on smaller scales decays faster and is re-created by turbulent cascading with new kinematical signature. This leaves no trace of the history of the fragment on small scales. Second, the kinematic structure of the core may also be affected by self-gravity: the change in kinematical properties appears to happen right at a Jeans length of about 3$\times$10$^4$ au which is also the core radius. Furthermore, the action of magnetic fields and density gradients can alter the behaviour of the turbulence. But the fact that today only a low level of turbulence is observed suppressing efficient grain growth hints towards an evolution which has seen stronger turbulence at former times. We have not made use of the other interesting peculiarity of the core: its strong depletion. Chemical model calculations for transient density fluctuations forming core-like density enhancements have been performed \citep{2005MNRAS.356..654G}. Under the assumption of a Gaussian core density profile whose amplitude varies as a time-dependent Gaussian, they probed the chemical evolution at various times and points in the core with a 221 species chemical code. It was found that the freeze-out in the high-density regions maintains for some time after the strongest compression has passed until the extinction drops below a critical value and the molecules are evaporated back into the gas phase. This result would be consistent with finding a core in the process of decreasing density and yet showing substantial depletion. While running a chemistry model for a core with the properties of L1506C and a given evolutionary sequence for the density structure is possible, comparison with the line data would also require to run a line transfer model. The parameter space to be covered will need careful exploration and this advanced effort is beyond the scope of this paper, but an interesting perspective for future work on this remarkable core. The presented results and both scenarios are in agreement with L1506C being a pre-stellar core in the making as proposed by Pea10: despite its turbulent and dense history the low-density core appears to be gravitational bound and showing infall motions. But within the picture of grain growth in cores, the presence of coreshine and depletion might indicate that it is in the re-making.	14	3	1403.3650	Context. With the initial steps of the star formation process in the densest part of the interstellar medium (ISM) still under debate, much attention is paid to the formation and evolution of pre-stellar cores. The recently discovered coreshine effect can aid in exploring the core properties and in probing the large grain population of the ISM. <BR /> Aims: We discuss the implications of the coreshine detected from the molecular cloud core L1506C in the Taurus filament for the history of the core and the existence of a primitive ISM component of large grains becoming visible in cores. <BR /> Methods: The coreshine surface brightness of L1506C is determined from Spitzer IRAC images at 3.6 μm. We perform grain growth calculations to estimate the grain size distribution in model cores similar in gas density, radius, and turbulent velocity to L1506C. Scattered light intensities at 3.6 μm are calculated for a variety of MRN and grain growth distributions using the DIRBE 3.5 μm all-sky map as external interstellar radiation field, and are compared to the observed coreshine surface brightness. <BR /> Results: For a core with the overall physical properties of L1506C, no detectable coreshine is predicted with a size distribution following the shape and size limits of an MRN distribution. Extending the distribution to grain radii of about 0.65 μm allows to reproduce the observed surface brightness level in scattered light. Assuming the properties of L1506C to be preserved, models for the growth of grains in cores do not yield sufficient scattered light to account for the coreshine within the lifetime of the Taurus complex. Only increasing the core density and the turbulence amplifies the scattered light intensity to a level consistent with the observed coreshine brightness. <BR /> Conclusions: The coreshine observed from L1506C requires the presence of grains with sizes exceeding the common MRN distribution. The grains could be part of primitive omni-present large grain population becoming visible in the densest part of the ISM, could grow under the turbulent dense conditions of former cores, or in L1506C itself. In the later case, L1506C must have passed through a period of larger density and stronger turbulence. This would be consistent with the surprisingly strong depletion usually attributed to high column densities, and with the large-scale outward motion of the core envelope observed today.	false	[ "model cores", "large grains", "cores", "former cores", "pre-stellar cores", "grain radii", "grain growth calculations", "grains", "Scattered light intensities", "larger density", "scattered light", "sufficient scattered light", "L1506C. Scattered", "L1506C", "MRN and grain growth distributions", "external interstellar radiation field", "size limits", "the grain size distribution", "primitive omni-present large grain population", "gas density" ]	11.403331	10.835609	-1
482976	[ "Christiaens, V.", "Casassus, S.", "Perez, S.", "van der Plas, G.", "Ménard, F." ]	2014ApJ...785L..12C	[ "Spiral Arms in the Disk of HD 142527 from CO Emission Lines with ALMA" ]	91	[ "Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile", "Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile", "Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile", "Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile", "UMI-FCA, CNRS/INSU, France (UMI 3386) and Departamento de Astronomía, Universidad de Chile, Santiago, Chile" ]	[ "2014ApJ...792...68S", "2014ApJ...792L..25V", "2014MNRAS.444.1919D", "2015A&A...578L...6B", "2015ASPC..499..273C", "2015Ap&SS.355..253G", "2015ApJ...798...85P", "2015ApJ...798L..44M", "2015ApJ...805...38D", "2015ApJ...807...10B", "2015ApJ...811...92C", "2015ApJ...812..126C", "2015ApJ...815L..21F", "2015MNRAS.448.3806L", "2015MNRAS.451..974D", "2015MNRAS.451.1147J", "2015MNRAS.452.3689E", "2015arXiv151002729M", "2016A&A...593A..11M", "2016A&A...596A..91H", "2016Ap.....59..449D", "2016ApJ...823L...8M", "2016ApJ...831..122R", "2016ApJ...832..178V", "2016IAUS..314..128C", "2016MNRAS.458..306H", "2016PASA...33...13C", "2016Sci...353.1519P", "2017A&A...597A..42B", "2017A&A...599A..86H", "2017A&A...604A..88W", "2017A&A...607A.114W", "2017AJ....154...33A", "2017ASSL..445..253D", "2017ApJ...840...32T", "2017ApJ...840...60B", "2017ApJ...850...52P", "2017ApJ...851...19B", "2017AstL...43..106D", "2017MNRAS.471..317R", "2017PASJ...69...97K", "2018A&A...617A..37C", "2018ApJ...852..122H", "2018ApJ...853..162B", "2018ApJ...854...84A", "2018ApJ...860..124D", "2018ApJ...864...81O", "2018ApJ...869L..41A", "2018ApJ...869L..43H", "2018ApJ...869L..44K", "2018MNRAS.475L..35M", "2018MNRAS.476.5115R", "2018MNRAS.477.1270P", "2018haex.bookE.165K", "2019A&A...622A..43C", "2019A&A...622A.147K", "2019A&A...624A..33V", "2019A&A...628A..20D", "2019ApJ...884L..56T", "2019MNRAS.482.3989R", "2019MNRAS.483.4114C", "2019MNRAS.489.2204P", "2020A&A...635L...9P", "2020A&A...637L...5B", "2020A&A...641A.128R", "2020ApJ...889L..24P", "2020ApJ...898..140H", "2020ApJ...900..135U", "2020ApJ...901...71S", "2020ApJ...904..125S", "2020ApJ...905...89Y", "2020ApJ...905..120B", "2020MNRAS.491..504C", "2020MNRAS.491.1335R", "2021A&A...650A..59B", "2021AJ....161...33V", "2021AJ....161..239F", "2021ApJ...912...56B", "2021ApJ...915..131L", "2021ApJ...923..128M", "2021ApJS..257...18T", "2021ApJS..257...19H", "2021MNRAS.504..782G", "2021MNRAS.507.3789C", "2023A&A...670L...8G", "2023ASPC..534..423B", "2023ASPC..534..645P", "2023ApJ...958...98F", "2024A&A...682A.101S", "2024A&A...683A...6N", "2024MNRAS.530.4802S" ]	[ "astronomy" ]	11	[ "planet-disk interactions", "protoplanetary disks", "stars: individual: HD 142527", "Astrophysics - Earth and Planetary Astrophysics" ]	[ "1964ApJ...139.1217T", "1979ApJ...233..857G", "1998ApJ...503..923B", "1999AJ....117..354D", "2001AJ....122.3396G", "2001ApJ...553..174G", "2001MNRAS.323..402L", "2003AJ....126..385C", "2004A&A...414.1153A", "2004A&A...421.1075D", "2004A&A...426..151A", "2004ApJ...605L..53F", "2005AJ....129.2481Q", "2005ApJ...622L.133C", "2006ApJ...636L.153F", "2006ApJ...644L.133F", "2006ApJ...645.1297L", "2008Ap&SS.313..101O", "2009A&A...493L..49H", "2010A&A...523A..25B", "2011A&A...528A..91V", "2011ApJ...738..131D", "2011ApJ...739..102K", "2011ApJ...740...84Q", "2012A&A...546A..24R", "2012A&A...547A..84T", "2012ApJ...748L..22M", "2012ApJ...753L..38B", "2012ApJ...754L..31C", "2012ApJ...755....6Z", "2013A&A...553A..64C", "2013A&A...556A.123C", "2013A&A...557A.133D", "2013A&A...560A..20B", "2013ApJ...762...48G", "2013ApJ...768..143Z", "2013MNRAS.430.1392R", "2013Natur.493..191C", "2013PASJ...65L..14F", "2014ApJ...781...87A", "2014ApJ...781L..30C" ]	[ "10.1088/2041-8205/785/1/L12", "10.48550/arXiv.1403.1463" ]	1403	1403.1463_arXiv.txt	The young circumstellar disks that host a large annular gap or central cavity, the so-called \emph{transition disks} (TDs), % could be crucial in the context of planetary formation as their morphologies may result from dynamical clearing due to planet(s) \citep{Dodson-Robinson2011, Zhu2012b, Zhu2013a}, rather than dust grain growth or photo-evaporation alone \citep{Rosotti2013}. Asymmetric features, warps or spirals in TDs can also evidence the presence of stellar or planetary companions \citep[e.g.][]{GoldreichTremaine1979}. Spiral features have mostly been detected in optical or near-infrared (NIR) wavelengths \citep[e.g.][]{Grady2001, Clampin2003, Fukagawa2004, Muto2012, Grady2013}, and in the disk of AB~Aur at radio wavelength \citep{Corder2005, Lin2006}, although in counter-rotation with the disk, thus probably stemming from a late envelope infall \citep{Tang2012}. The disk around the Herbig Fe star \objectname{HD~142527} constitutes an outstanding example of a nearly face-on TD \citep[$i$$\sim$28$\pm$3\degr;][]{Perez2014}, showing both a large asymmetric dust depleted gap and a spiral arm observed in NIR scattered light \citep{Fukagawa2006}. Its mass and age were estimated to respectively $\sim$1.9-2.2 M$_{\odot}$ and 1-12 Myr old \citep{Fukagawa2006, Verhoeff2011}, for a distance of 145$\pm$15 pc \citep{DeZeeuw1999, Acke2004}. \citet{Fujiwara2006} concluded that the disk's rotation axis was inclined so that the east side was the far side, with a position angle (PA) of $\sim$-20\degr. With K-band imaging, \citet{Casassus2012} reported the presence of smaller spiral structures at the outer edge of the gap, confirmed by \citet{Rameau2012} and in NIR polarized intensity studies \citep[][the latter work identifies two new spirals]{Canovas2013, Avenhaus2014}. ALMA provided substantial insight into the structure of HD~142527. \citet{Casassus2013} found gap-crossing HCO$^+$(4-3) filaments, diffuse CO(3-2) inside the cavity, and confirmed the horseshoe morphology of the dust continuum reported by \citet{Ohashi2008}, which they interpreted as a dust trap. \citet{Fukagawa2013} reported on another ALMA dataset in $^{13}$CO J=3-2, C$^{18}$O J=3-2 and underlying continua. Here we focus on the outskirts of the \objectname{HD 142527} disk % and report on the discovery of several CO spiral arms.	\label{sec:discussion} % Several scenarios have been proposed for the occurrence of spirals in TDs \citep[see e.g.][for a summary]{Boccaletti2013}. HD~142527, with its well-defined tightly coiled radio spirals in Keplerian rotation, provides an interesting case study suggesting a different origin than the CO spirals in the disk of AB~Aur. S1 could be fit with equations~\eqref{Eq:Muto} and \eqref{Eq:Kim} assuming an embedded companion (section \ref{sec:ModellingSpirals}), although its location on the model spiral could not be precisely constrained. % An object has been discovered by \citet{Biller2012} and re-detected in H$\alpha$ by \citet{CloseFollette2014} at $\sim$12~AU. Extending S1 inwards does not allow us to confirm a possible relationship. Nevertheless, the clearing of a gap as large as 140~AU should involve several planets \citep[e.g.][]{Dodson-Robinson2011}. For S2 and S3, both the \emph{relatively} poor least-square fit with equations \eqref{Eq:Muto} and \eqref{Eq:Kim} and their very large scales argue in favor of a different cause.% Simulations in the context of protoplanetary disks have shown that two-armed spirals, such as S2 and S3, could be induced by stellar encounters \citep[e.g.][]{Larwood2001,Quillen2005}. Similarly, simulations of \citet{Augereau2004} and \citet{Quillen2005} reproduce the large scale ($\sim$325 AU) tightly wound spiral observed in the disk of \objectname{HD~141569~A}, by tidal induction from one of its M-dwarf companions (external to the disk). Tidally induced spiral structures are transient and vanish within a few orbits, which implies either a very recent encounter or a bound companion external to the disk periodically exciting spirals. For HD~142527 no such object has been detected. Gravitational instability (G.I.) is an alternative scenario able to create a grand-design two-armed spiral structure \citep[see e.g.][]{Boss1998}. The stability of a disk is characterized by Toomre's parameter, $Q$ \citep{Toomre1964}. \citet{Fukagawa2013} computed an upper limit of 1-2 for $Q$ in the dust continuum horseshoe, so that depending on the gas-to-dust ratio, fragmentation of either the gas or the dust or both components due to G.I. may occur. Approximating $Q \approx \frac{M_{\star}}{M_d} h_S$ \citep{Gammie2001}, with the mass of HD 142527 $M_{\star}$ set to 2 $M_{\odot}$, the mass of the disk $M_d$ set to 0.1 $M_{\odot}$ \citep{Verhoeff2011} and the aspect ratio at both spirals $h_S$$\sim$0.1 (section \ref{sec:DescriptionSpirals}), we find $Q$$\sim$2.0 similarly to \citet{Fukagawa2013}. This suggests that the outer disk in general, not just the horse-shoe, is stable but close to the instability regime.	14	3	1403.1463	In view of both the size of its gap and the previously reported asymmetries and near-infrared spiral arms, the transition disk of the Herbig Fe star HD 142527 constitutes a remarkable case study. This paper focuses on the morphology of the outer disk through ALMA observations of <SUP>12</SUP>CO J = 2-1, <SUP>12</SUP>CO J = 3-2, and <SUP>13</SUP>CO J = 2-1. Both <SUP>12</SUP>CO J = 2-1 and <SUP>12</SUP>CO J = 3-2 show spiral features of different sizes. The innermost spiral arm (S1) is a radio counterpart of the first near-infrared spiral observed by Fukagawa, but it is shifted radially outward. However, the most conspicuous CO spiral arm (S2) lies at the outskirts of the disk and has not been detected before. It corresponds to a cold density structure, with both brightness and excitation temperatures of order 13±2 K and conspicuous in the <SUP>12</SUP>CO J = 2-1 peak-intensity map, but faint in <SUP>12</SUP>CO J = 3-2. There is also a faint counterarm (S3), at a point-symmetric location of S2 with respect to the star. These three spirals are modeled separately with two different formulae that approximate the loci of density maxima in acoustic waves due to embedded planets. S1 could be fit relatively well with these formulae, compared to S2 and S3. Alternative scenarios such as gravitational instability or external tidal interaction are discussed. The impact of channelization on spectrally and spatially resolved peak intensity maps is also briefly addressed.	false	[ "CO J", "CO", "SUP>12</SUP", "embedded planets", ">CO J", "different sizes", "a remarkable case study", "density maxima", "external tidal interaction", "acoustic waves", "Herbig Fe", "S2", "S3", "ALMA observations", "the most conspicuous CO spiral arm", "respect", "The innermost spiral arm", "gravitational instability", "acoustic", "maxima" ]	9.601627	12.743021	149
871039	[ "Zaninetti, L." ]	2014RMxAA..50....7Z	[ "A near infrared test for two recent luminosity functions for galaxies" ]	3	[ "Dipartimento di Fisica, Università degli Studi di Torino, Italy" ]	[ "2018OAst...27..335Z", "2019Galax...7...61Z", "2019IJAA....9..393Z" ]	[ "astronomy" ]	4	[ "galaxies: clusters: general", "galaxies: luminosity function", "mass function", "galaxies: statistics", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]	[ "1929PNAS...15..168H", "1976ApJ...203..297S", "1996catp.book.....P", "2001MNRAS.326..255C", "2002thas.book.....P", "2004BrJPh..34..455A", "2011ApJ...734...40W", "2011MNRAS.414.1903G", "2012ApJ...754..131K", "2012ApJ...758...24F", "2012ApJS..199...26H", "2012RMxAA..48..209Z" ]	[ "10.48550/arXiv.1403.7073" ]	1403	1403.7073_arXiv.txt	The release of the 2MASS Redshift Survey (2MRS), with it's 44599 galaxies having $K_s < 11.75$ allows to make tests on the radial number of galaxies because we have a small zone-of-avoidance , see Figure 1 in \citet{Huchra2012}. The number of galaxies as function of the redshift , $z$, is strictly related to the chosen luminosity function for galaxies (LF). The most used LF is the Schechter function , introduced by \citet{schechter} , but also two recent LFs, the generalized gamma with four parameters, see \citet{Zaninetti2010f}, and the modified Schechter LF , see \citet{Alcaniz2004}, can model the LF for galaxies. We now outline some topic issues in which the LF plays a relevant role : determination of the parameters at different $z$ , see \citet{Goto2011}, evaluation of the parameters taking account of the star formation (SF) and presence of active nuclei, see \citet{Wu2011}, determination of the normalization in the Near Infrared (NIR) as function of $z$, see \citet{Keenan2012}. In this paper Section \ref{seclum} first reviews two recent LFs for galaxies and then derives the free parameters in the near infrared band. Once the basic parameters of the recent LFs are derived we make a comparison between observed radial distribution in the number of galaxies and theoretical predictions, see Section \ref{secradial}. A simulation of the near infrared all sky survey is reported in Section \ref{simulation}.	The most used LF in the infrared band is the Schechter LF, see eqn. (\ref{lfstandard}). We have here analyzed two other LFs : the generalized gamma, see eqn. (\ref{lfgamma}) and the modified Schechter LF, eqn. (\ref{lfbrazilian}). They perform as well as the Schechter LF and the test in the $K_s$ band assigns to the modified Schechter LF the lowest value of the $\chi^2_{red}$, see Table \ref{chi2value}. The radial distribution in the number of galaxies of 2MRS can be another test and Figure \ref{maximum_flux_brazilian} and Figure \ref{maximum_flux_gamma4} report the standard theoretical curve as well the two new theoretical predictions. In this case the smaller $\chi^2$ is given by the theoretical curve which involves the modified Schechter LF. A particular attention has been given to the position of the maximum in the number of galaxies that is here expressed as function of the theoretical parameter $z_{crit}$ or the two observable parameters $f$ and $m$. A simulation in Mercator projection of the spatial distribution of galaxies having redshift corresponding to the photometric maximum is presented, see Figure \ref{sphere_voronoi_galaxies} .	14	3	1403.7073	Two recent luminosity function (LF) for galaxies are reviewed and the parameters which characterize the near infrared are fixed. A first LF is a modified Schechter LF with four parameters. The second LF is derived from the generalized gamma distribution and has four parameters. The formulas which give the number of galaxies as function of the redshift are reviewed and a special attention is given to the position of the photometric maximum which is expressed as function of a critical parameter or the flux of radiation or the apparent magnitude. A simulation of the 2MASS Redshift Survey is given in the framework of the non Poissonian Voronoi Tessellation.	false	[ "function", "Schechter LF", "radiation", "LF", "the non Poissonian Voronoi Tessellation", "a critical parameter", "the apparent magnitude", "galaxies", "Redshift Survey", "four parameters", "the parameters", "the near infrared", "a modified Schechter LF", "Two recent luminosity function", "the generalized gamma distribution", "the photometric maximum", "a special attention", "A first LF", "The second LF", "the 2MASS Redshift Survey" ]	12.327982	5.942861	189
437376	[ "Rauch, T.", "Werner, K.", "Quinet, P.", "Kruk, J. W." ]	2014A&A...564A..41R	[ "Stellar laboratories. II. New Zn iv and Zn v oscillator strengths and their validation in the hot white dwarfs G191-B2B and RE 0503-289" ]	26	[ "Institute for Astronomy and Astrophysics, Kepler Center for Astro and Particle Physics, Eberhard Karls University, Sand 1, 72076, Tübingen, Germany", "Institute for Astronomy and Astrophysics, Kepler Center for Astro and Particle Physics, Eberhard Karls University, Sand 1, 72076, Tübingen, Germany", "Astrophysique et Spectroscopie, Université de Mons - UMONS, 7000, Mons, Belgium; IPNAS, Université de Liège, Sart Tilman, 4000, Liège, Belgium", "NASA Goddard Space Flight Center, Greenbelt, MD, 20771, USA" ]	[ "2014A&A...566A..10R", "2015A&A...574A..29W", "2015A&A...577A...6R", "2015A&A...577A..88R", "2015A&A...578A.125H", "2015A&A...582A..94W", "2015ASPC..493....3C", "2015ASPC..493...27W", "2015PhDT.......250R", "2016A&A...587A..39R", "2016A&A...590A.128R", "2017A&A...598A.135H", "2017A&A...599A.142R", "2017A&A...606A.105R", "2017ASPC..509..183R", "2017ASPC..509..189H", "2017ASPC..509..235L", "2017CaJPh..95..790Q", "2018A&A...612A..62H", "2018MNRAS.480.1211R", "2019MNRAS.489.1054L", "2020A&A...637A...4R", "2020MNRAS.492..528L", "2021A&A...647A.184R", "2023MNRAS.518..368C", "2024Ap&SS.369...43B" ]	[ "astronomy", "physics" ]	5	[ "atomic data", "line: identification", "stars: abundances", "stars: individual: G191-B2B", "virtual observatory tools", "stars: individual: RE 0503-289", "Astrophysics - Solar and Stellar Astrophysics", "Physics - Atomic Physics" ]	[ "1981tass.book.....C", "1994ApJ...425L.105H", "1995JPCRD..24.1803S", "1996A&A...314..217D", "2003ASPC..288...31W", "2003ASPC..288..103R", "2009ARA&A..47..481A", "2012A&A...546A..55R", "2012ApJ...753L...7W", "2013A&A...560A.106R" ]	[ "10.1051/0004-6361/201423491", "10.48550/arXiv.1403.2183" ]	1403	1403.2183_arXiv.txt	\label{sect:intro} In a recent spectral analysis of the hydrogen-rich DA-type white dwarf \gb, \citet{rauchetal2013} identified and reproduced stellar lines of C, N, O, Al, Si, O, P, S, Fe, Ni, Ge, and Sn. In addition, they identified 21 \ion{Zn}{iv} lines. The determined Zn abundance (logarithmic mass fraction of $-4.89$, 7.5 $\times$ solar) was uncertain because the unknown \ion{Zn}{iv} oscillator strengths were approximated by values of the isoelectronic \ion{Ge}{vi} taken from \citet{rauchetal2012}. In this paper, we introduce new oscillator strengths for \ion{Zn}{iv} and \ion{Zn}{v} (Sect.\,\ref{sect:zntrans}). Then, we describe briefly our observations (Sect.\,\ref{sect:observation}), our analysis strategy (Sect.\,\ref{sect:models}), and revisit \gb to perform a precise determination of its Zn abundance (Sect.\,\ref{sect:model}). The white dwarf \re is hotter than \gb and its trans-iron element abundances are strongly oversolar \citep{werneretal2012, rauchetal2013} and, thus, it appears promising to identify Zn lines. In Sect.\,\ref{sect:re}, we describe our search for these and the determination of its Zn abundance. We summarize our results and conclude in Sect.\,\ref{sect:results}.	\label{sect:results} The identified \ion{Zn}{iv} and \ion{Zn}{v} lines in the high-resolution UV spectra of \gb and \re are well reproduced with our newly calculated oscillator strengths by our NLTE model-atmosphere calculations. We determined photospheric abundances of $\log\,\mathrm{Zn} = -5.52 \pm 0.2$ (mass fraction, $1.9 - 4.8\,\times\,10^{-6}$, 1.1 -- 2.8 times the solar abundance) and $\log\,\mathrm{Zn} = -3.57 \pm 0.2$ ($1.7 - 4.3\,\times\,10^{-4}$, 98 -- 248 times solar) for the DA-type white dwarf \gb and the DO-type white dwarf \re, respectively. The highly supersolar Zn abundance is in line with the high abundances of trans-iron elements Ge \citep[650\,$\times$ solar,][]{rauchetal2012}, Kr (450\,$\times$ solar), Xe \citep[3800\,$\times$ solar,][]{werneretal2012} in \re. The identification of new lines due to trans-iron elements, e.g., Ga, Ge, As, Se, Kr, Mo, Sn, Te, I, and Xe \citep{werneretal2012} and Zn \citep[and in this paper]{rauchetal2013} in \gb and \re promises to help enhance the understanding of extremely metal-rich white dwarf photospheres and their relation to AGB and post-AGB stellar evolution. Their reproduction, i.e., the precise abundance determination, e.g., of Kr and Xe \citep{werneretal2012}, Ge and Sn \citep{rauchetal2012}, and Zn (this paper) is strongly dependent on the available atomic data. This remains a challenge for atomic and theoretical physicists.	14	3	1403.2183	Context. For the spectral analysis of high-resolution and high-signal-to-noise (S/N) spectra of hot stars, state-of-the-art non-local thermodynamic equilibrium (NLTE) model atmospheres are mandatory. These are strongly dependent on the reliability of the atomic data that is used for their calculation. In a recent analysis of the ultraviolet (UV) spectrum of the DA-type white dwarf <ASTROBJ>G191-B2B</ASTROBJ>, 21 Zn iv lines were newly identified. Because of the lack of Zn iv data, transition probabilities of the isoelectronic Ge vi were adapted for a first, coarse determination of the photospheric Zn abundance. <BR /> Aims: Reliable Zn iv and Zn v oscillator strengths are used to improve the Zn abundance determination and to identify more Zn lines in the spectra of <ASTROBJ>G191-B2B</ASTROBJ> and the DO-type white dwarf <ASTROBJ>RE 0503-289</ASTROBJ>. <BR /> Methods: We performed new calculations of Zn iv and Zn v oscillator strengths to consider their radiative and collisional bound-bound transitions in detail in our NLTE stellar-atmosphere models for the analysis of the Zn iv - v spectrum exhibited in high-resolution and high-S/N UV observations of <ASTROBJ>G191-B2B</ASTROBJ> and <ASTROBJ>RE 0503-289</ASTROBJ>. <BR /> Results: In the UV spectrum of <ASTROBJ>G191-B2B</ASTROBJ>, we identify 31 Zn iv and 16 Zn v lines. Most of these are identified for the first time in any star. We can reproduce well almost all of them at log Zn = -5.52 ± 0.2 (mass fraction, about 1.7 times solar). In particular, the Zn iv / Zn v ionization equilibrium, which is a very sensitive T<SUB>eff</SUB> indicator, is well reproduced with the previously determined and log g = 7.60 ± 0.05. In the spectrum of <ASTROBJ>RE 0503-289</ASTROBJ>, we identified 128 Zn v lines for the first time and determined log Zn = -3.57 ± 0.2 (155 times solar). <BR /> Conclusions: Reliable measurements and calculations of atomic data are a pre-requisite for stellar-atmosphere modeling. Observed Zn iv and Zn v line profiles in two white dwarf (<ASTROBJ>G191-B2B</ASTROBJ> and <ASTROBJ>RE 0503-289</ASTROBJ>) ultraviolet spectra were well reproduced with our newly calculated oscillator strengths. This allowed us to determine the photospheric Zn abundance of these two stars precisely. <P />Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26666.Based on observations made with the NASA-CNES-CSA Far Ultraviolet Spectroscopic Explorer.Tables 1 and 2 are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/564/A41">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/564/A41</A>	false	[ "Zn v oscillator strengths", "Zn iv", "Zn", "Zn iv data", "Reliable Zn", "Reliable Zn iv and Zn v oscillator strengths", "ASTROBJ", "more Zn lines", "ESA Hubble Space Telescope", "and Zn v line profiles", "determined log", "hot stars", "128 Zn v lines", "16 Zn v lines", "the Zn abundance determination", "NASA", "the photospheric Zn abundance", "ultraviolet spectra", "NLTE", "RE" ]	8.060555	10.337361	-1
800163	[ "Castel-Branco, Nuno", "Páramos, Jorge", "March, Riccardo" ]	2014PhLB..735...25C	[ "Perturbation of the metric around a spherical body from a nonminimal coupling between matter and curvature" ]	19	[ "Instituto Superior Técnico, Universidade de Lisboa, Av. Rovisco Pais 1, 1049-001 Lisboa, Portugal", "Centro de Física do Porto and Departamento de Física e Astronomia, Faculdade de Ciências, Universidade do Porto, Rua do Campo Alegre 687, 4169-007, Porto, Portugal", "Istituto per le Applicazioni del Calcolo, CNR, Via dei Taurini 19, 00185 Roma, Italy" ]	[ "2013arXiv1306.1177B", "2014Galax...2..410H", "2014arXiv1407.2751C", "2015GReGr..47.1835B", "2016Ap&SS.361..365H", "2016arXiv161106606A", "2017AdSpR..59..645D", "2017AdSpR..60.1300C", "2017PhRvD..95b4017M", "2018CQGra..35b5002S", "2018EPJC...78..957A", "2019PhRvD.100d2002M", "2019arXiv190307059M", "2021PhRvD.104b4052W", "2022NatSR..12.3981W", "2022PDU....3701090S", "2022PhRvD.105b4020F", "2022ResPh..3805594G", "2024JCAP...05..115W" ]	[ "astronomy", "physics" ]	6	[ "f (R ) theories", "Nonminimal coupling", "Solar System", "f (R) theories", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]	[ "1980PhLB...91...99S", "1980PhRvD..22..343D", "1990CQGra...7.1023S", "1992CQGra...9..895V", "1994CQGra..11..269W", "2003ARNPS..53...77A", "2004PhRvD..69d4026K", "2004PhRvD..70d3528C", "2004PhRvL..93q1104K", "2005gr.qc.....9114M", "2006LRR.....9....3W", "2006MNRAS.371..626S", "2006astro.ph.11816W", "2006gr.qc.....2016B", "2007arXiv0706.1024B", "2008CQGra..25x5017B", "2008IJMPA..23.4817B", "2008PhRvD..77b4041C", "2008PhRvD..78f4036B", "2009ApJS..180..330K", "2010JCAP...03..009B", "2010PhRvD..81j4046B", "2010PhRvL.105w1103L", "2011PhRvD..83d4010B", "2011PhRvD..84f4022B", "2011PhRvL.106v1101E", "2012PhRvD..85b4012B", "2012PhRvD..86d4034B", "2012PhRvD..86j3007P", "2012arXiv1212.2177B", "2013JCAP...05..029B", "2013PhLA..377.2447P", "2013PhRvD..87d4045P", "2013PhRvD..87h1502O", "2013PhRvD..88f4019B", "2013PhRvD..88f4025P", "2013arXiv1306.1177B", "2014CQGra..31h5003I", "2014JMoSp.300...65S", "2014PhRvD..89d4012B" ]	[ "10.1016/j.physletb.2014.06.001", "10.48550/arXiv.1403.7251" ]	1403	1403.7251_arXiv.txt	Modern physics uses the concepts of dark matter and dark energy to advance an explanation for the astrophysical problem of the flattening of galactic rotation curves and the cosmological problem of the accelerated expansion of the universe, respectively. Dark energy, which is supposed to account for 74$\%$ of all the matter of the universe, has many theories on its basis, as the so-called "quintessence" models \cite{quintessence1,quintessence2,quintessence3} and the existence of scalar fields that account for both dark matter and dark energy \cite{scalarfield}. More recent approaches start from the idea of the incompleteness of the fundamental laws of General Relativity (GR), involving, for example, corrections to the Einstein-Hilbert action. Such theories involve a nonlinear correction to the geometry part of the action, being thus called $f(R)$ theories. In the last decade, work on $f(R)$ theories has been very profitable, as thoroughly discussed in Ref. \cite{felice}. These can be extended to also include a nonminimum coupling (NMC) between the scalar curvature and the matter Lagrangian density. Indeed, these NMC theories have many interesting features, as can be seen by several studies, such as the impact on stellar observables \cite{stelobserv}, the energy conditions \cite{energcondit}, the equivalence with multi-scalar-tensor theories \cite{multiscalar}, the possibility to account for galactic \cite{drkmattgal} and cluster \cite{drkmattclus} dark matter, cosmological perturbations \cite{cosmpertur}, a mechanism for mimicking a Cosmological Constant at astrophysical scales \cite{mimlambda}, post-inflationary reheating \cite{reheating} or the current accelerated expansion of the universe \cite{curraccel,Friedmann}, the dynamical impact of the choice of the Lagrangian density of matter \cite{dynimpac1,dynimpac2}, gravitational collapse \cite{gravcollapse}, its Newtonian limit \cite{newtlimit}, the existence of closed timelike curves \cite{closedtimecurve} and, the most recent one, a determination of solar system constraints to a cosmological NMC \cite{solar}. For other NMC gravity theories and their potential applications, see {\it e.g.} \cite{puetz1,puetz2,puetz3,obuk,iorio}. One of the first motivations that brought $f(R)$ theories into the physicists daily work was the Starobinsky inflation model, where $f(R)=R+R^2/(6m^2)$ was considered \cite{starobinsky,reheating}, with WMAP normalization of the CMB temperature anisotropies indicating that $m\sim 3 \times 10^{-6} M_P$, where $M_P$ is the Planck mass \cite{planckmass}. Without directly mentioning the Starobinsky inflation, Ref. \cite{naf} considers a quadratic $f(R)$ function and develops an expansion in powers of $ (1/c)$ of an asymptotically flat Minkowski metric, showing the presence of a Yukawa correction to the $tt$ component of the latter \cite{naf}. Following the equivalence between scalar-tensor and $f(R)$ theories \cite{analogy1,analogy2,analogy3}, this can be interpreted as due to the additional gravitational contribution of the massive degree of freedom embodied in a non-linear $f(R)$ function. In this work we follow a similar procedure of Ref. \cite{naf} where we instead consider a NMC model. We consider that the additional degree of freedom arising from a non-trivial $f(R)$ function is sufficiently massive so that its effects are not extremely long-ranged, and as such we can neglect the background cosmological setting --- this point (which shall be developed in the following) shows that this work is complementary to the recent study on the compatibility between cosmological and solar system dynamics of a NMC model \cite{solar}. In section II such a model is presented and in section III the solution of the linearized field equations is computed. We obtain the solutions for the perturbative potentials $\Psi(r)$ and $\Phi(r)$ of the metric, which contain a form factor specific of the Yukawa potential that is addressed in section IV. The $tt$ component of the metric yields the modified gravitational potential, which includes a Newtonian plus a Yukawa contribution. The comparison of these results with available experimental constrains is presented in section V. This section also addresses the radial potential through the constraints obtained to the geodetic precession values. Finally, conclusions are drawn.	As it has been shown, from the $tt$ component of the metric we identify a Newtonian potential plus a Yukawa perturbation: \begin{equation} \label{yukawa potential} U(r) = - \dfrac{G M_S}{r} \left( 1 + \alpha A(m,R_S) e^{-r/\lambda} \right), \end{equation} \noindent defining the characteristic length $\lambda = 1/m$ and the strength of the Yukawa addition \begin{equation} \label{alpha as the yukawa potential strength} \alpha = \dfrac{1}{3} - 4\xi , \end{equation} \noindent so that, if $\xi = 0$, we get the Yukawa strength for pure $f(R)$ theories, $\alpha = 1/3$ \cite{naf}; also, notice that a positive NMC (as assumed in Refs. \cite{stelobserv,gravcollapse}) yields $\alpha \leq 1/3$. Strikingly, a NMC with $\xi = 1/12$ cancels the Yukawa contribution. \subsection{The form factor $A(m,R_S)$} As defined before, the form factor is \begin{equation} \label{form factor A(m,R_S)} A(m,R_S) = \dfrac{4\pi}{m M_S} \int_0^{R_S} \sinh(mr) r\rho(r)dr. \end{equation} This dimensionless form factor was found by integrating the field equations of NMC gravity but it is not specific of the NMC gravity nor of $f(R)$ theories of this kind, but of any Yukawa model \cite{fischbach}, as it arises from the integral \begin{equation} \label{yukawa potential, integral} U_Y \left( \vec{x} \right) = - G \alpha \int_{B_{R_S}} \dfrac{\exp\left[-m \left\| \vec{x}-\vec{x}'\right\| \right]}{\left\| \vec{x}-\vec{x}'\right\|} \rho \left(\vec{x}'\right) d^3x', \end{equation} \noindent where $B_{R_S}$ is a sphere with radius $R_S$ and center at the origin. In the case of a spherically symmetric distribution of mass $\rho(r)$, evaluation of the integral \eqref{yukawa potential, integral} in spherical coordinates yields the Yukawa contribution to $U(r)$ in Eq. \eqref{yukawa potential}, so that $G\alpha$ can be interpreted as the strength of the Yukawa potential generated by a point source. The form factor can be evaluated in several ways, according to the function of mass density $\rho(r)$. Taking the limit of a point source, $r \to 0$ allows us to expand around $m r \ll 1$, so that $\sinh \left( m r \right) \approx m r [1 + (mr)^2/6]$ and \begin{equation} A(m,R_S) \approx \dfrac{4\pi}{M_S} \int^{R_S}_0 \left[1 + \dfrac{(mr)^2 }{6}\right] r^2 \rho(r) dr = 1 + \dfrac{2m^2\pi}{ 3M_S} \int^{R_S}_0 r^4 \rho(r) dr \sim 1. \end{equation} \noindent This can be verified explicitly by making all computations with a test mass density (such as a uniform profile) and, in the end, taking the limit $R_S \to 0$. Indeed, taking \begin{equation} \rho_0 = \dfrac{3M_S}{4\pi R_S^3} \end{equation} \noindent we obtain \begin{equation} \label{form factor A(m,R_S) constant} A(m,R_S) =3 \dfrac{mR_S \cosh (mR_S) - \sinh (mR_S)}{(m R_S)^3} , \end{equation} \noindent which admits the limiting cases \begin{eqnarray} \label{form factor A(m,R_S) constant limiting} A(m,R_S) &\approx & 1 + \dfrac{(mR_S)^2 }{ 10} \sim 1 ~~,~~mR_S \ll 1, \\ \nonumber A(m,R_S) &\approx & \dfrac{3}{2}\dfrac{e^{mR_S}}{(mR_S)^2} ~~,~~mR_S \gg 1. \end{eqnarray} If the central body is the Sun (with radius $R_\odot$), we may instead consider the more accurate density profile \cite{NASAprofile} (which obeys condition $\rho(R_\odot)=0$, while $(d\rho\slash dr)(R_\odot) \simeq 0$), \begin{equation}\label{NASA density profile} \rho(r) = \rho_0 \bigg[ 1 - 5.74 \left(\dfrac{r}{R_\odot}\right)+ 11.9\left(\dfrac{r}{R_\odot}\right)^2 - 10.5\left(\dfrac{r}{R_\odot}\right)^3 + 3.34\left(\dfrac{r}{R_\odot}\right)^4\bigg] , \end{equation} \noindent obtaining \begin{eqnarray} \label{form factor A(m,R_S) NASA} A(m,R_\odot) &=& x^{-7} [4.6 \times 10^4 x + 2.1 \times 10^3 x^3 + (2.7 \times 10^4 + 131 x^2) x \cosh x - \\ \nonumber && (7.3 \times 10^4 + 3.6 \times 10^3x^2 - 14.6 x^4) \sinh x ], \end{eqnarray} \noindent (with $x= mR_\odot$, for brevity), with the limiting cases \begin{eqnarray} \label{form factor A(m,R_S) NASA limiting} A(m,R_\odot) &\approx & 1 + 6 \times 10^{-2} (m R_\odot)^2 \sim 1 ~~,~~mR_\odot \ll 1, \nonumber \\ A(m,R_\odot) &\approx & \ 7.3 \dfrac{e^{mR_\odot}}{(mR_\odot)^3 }~~,~~mR_\odot \gg 1. \end{eqnarray} Both forms for $A(m,R_\odot)$ are plotted in Fig. \ref{fig:formfactors}, showing that it grows with $m$. Although, for values of the lengthscale $\lambda \ll R_\odot$, this effectively boosts the form factor, the contribution from the Yukawa term in Eq. (\ref{yukawa potential}) is nevertheless suppressed by the factor $\exp(-r/\lambda)$. \subsection{PPN Parameters} Similarly to the present work, the Parameterized Post-Newtonian formalism posits an expansion of the metric elements and other quantities (energy-momentum tensor, equations of motion, {\it etc}.) in powers of $1/c^2$ \cite{Will}; the eponymous PPN metric reads, for a spherical central body, \begin{equation} \label{metricPPN} ds^2 = - \left[ 1 - 2\dfrac{GM_S }{c^2 r} + 2 \beta \left(\dfrac{ GM_S }{c^2 r}\right)^2 \right]c^2 dt^2 + \left( 1 + 2 \gamma \dfrac{GM_S }{c^2 r} \right)~\left( dr^2 + r^2 d\Omega^2 \right) ~~,\end{equation} \noindent where $\beta$ and $\gamma$ are two PPN parameters, which measure the amount of non-linearity affecting the superposition law for gravity and the spatial curvature per unit mass, respectively; GR is signalled by $\beta = \gamma =1$. Other PPN parameters also appear in a more evolved version of the metric above, signalling violation of momentum conservation, existence of a privileged reference frame, amongst others deviations from GR. Clearly, such a formalism is incompatible with the presence of a Yukawa term in the gravitational potential, since the latter cannot be expanded in powers of $1/r$; furthermore, the discussion after Eq. (\ref{curvature solution, outside star}) highlights that, in the limit $m\to 0$, we must consider the background cosmological curvature and cannot assume the asymptotically flat {\it Ansatz} (\ref{metric}) for the metric: this was performed in Ref. \cite{solar}, as already mentioned. Nevertheless, we can consider what happens if the condition $mr \ll 1$ is valid throughout the region of interest ({\it e.g.} the Solar System), for consistency: in this case, the metric (\ref{metric}) with the solutions (\ref{psi solution}) and (\ref{phi solution}) is well approximated by \begin{equation} \label{metriclight} ds^2 = -\left[1 - \dfrac{2G M_S}{c^2 r} \left(\dfrac{4}{3} - 4\xi \right) \right]c^2dt^2+ \left[1+\dfrac{2G M_S}{c^2 r} \left( \dfrac{2}{3}+ 4\xi \right)\right] dr^2 + r^2 d\Omega^2, \end{equation} \noindent which, upon comparison with Eq. (\ref{metricPPN}) (or, mathematically, the adequate constant rescaling of both time and radial coordinates), yields \begin{equation}\label{gamma} \gamma= \dfrac{1 }{ 2} \dfrac{1 + 6\xi }{ 1 - 3\xi}. \end{equation} In the absence of a NMC, $\xi = 0 $, this yields $\gamma = 1/2$, a strong departure from GR that is disallowed by current experimental bounds, $\gamma = 1 + (2.1 \pm 2.3) \times 10^{-5}$ \cite{status}. This apparent disagreement between $f(R)$ theories and observations was noted early on (as discussed {\it e.g.} in Refs. \cite{PPNfR1,PPNfR2,CSE}), and can be avoided if the additional degree of freedom arising from a non-linear $f(R)$ function is massive enough. The expression above appears to show that a NMC allows $f(R)$ theories to remain compatible with observations, as long as $\xi=1/12$ --- which is just a restatement of the previously obtained result. Again, the path towards obtaining the $\gamma$ PPN parameter depicted above is presented for illustration only, as it relies on an approximation of a Yukawa perturbation and disregards the fact that, in the limit $mr \ll 1$, the background cosmological dynamics cannot be neglected. As such, no conclusions can be drawn from comparison with the experimental bound on $\gamma$ mentioned above. In this work we have computed the effect of a NMC model, specified by \eqref{f(R) equations}, in a perturbed weak-field Schwarzschild metric, as depicted in Eq. (\ref{psi solution}) and (\ref{phi solution}). In the weak-field limit, this translates into a Yukawa perturbation to the usual Newtonian potential, with characteristic range and coupling strength \begin{equation} \lambda = \dfrac{1}{m}~~~~,~~~~\alpha = \left( \dfrac{1}{3} - 4\xi \right) = \dfrac{1}{3} \left[ 1 - \left( \dfrac{m}{M} \right)^2 \right]~~. \end{equation} \noindent This result is quite natural and can be interpreted straightforwardly: a minimally coupled $f(R)$ theory introduces a new massive degree of freedom (as hinted by the equivalence with a scalar-tensor theory \cite{analogy1,analogy2,analogy3}), leading to a Yukawa contribution with characteristic lengthscale $\lambda = 1/m$ and coupling strength $\alpha = 1/3$. The introduction of a NMC has no dynamical effect in the vacuum, as there is no matter to couple the scalar curvature to: as a result, we do not expect any modification in the range of this Yukawa addition; conversely, a NMC has an impact on the description of the interior of the central body (as illustrated in Refs.\cite{stelobserv,mimlambda,dynimpac2}), leading to a correction to the latter's coupling strength (which has a negative sign since $\LL_m= -\rho$). Using the available experimental constraints, we find that, for $10^{-22}~{\rm eV} < m < 1~{\rm meV}$ ({\it i.e.} the range $ 10^{-4}~{\rm m} < \lambda < 10^{16}~{\rm m}$), where $\|\alpha\| \ll 1$, we must have $\xi \sim 1/12$ or, equivalently, that both mass scales $m$ and $M$ of the non-trivial functions $f^1(R)$ and $f^2(R)$ must be extremely close. If this is the case, the latter relation is not interpreted as an undesirable fine-tuning, but instead is suggestive of a common origin for both non-trivial functions $f^1(R)$ and $f^2(R)$, in line with the argument stating that the model \eqref{model} should arise as a low energy phenomenological approximation to a yet unknown fundamental theory of gravity. Conversely, for values of $m$ (or $\lambda$) away from the range mentioned above the Yukawa coupling strength $\alpha$ can be much larger than unity, so that $\xi $ can assume any value and the mass scales $m$ and $M$ can differ considerably. In particular, the Starobinsky inflationary model, which requires the much heavier mass scale $m \approx 3 \times 10^{13}$ GeV $\sim 10^{-6} M_P$, manifests itself at a lengthscale $\lambda \sim 10^{-29}~{\rm m}$. This implies that the generalized preheating scenario posited in Ref. \cite{reheating}, which requires $1<\xi<10^4$, is thus completely allowed by experiment and unconstrained by this work. By computing the perturbation induced on geodetic precession, we have found that no significant new constraint arises, as this is already included in the existing Yukawa exclusion plot. Furthermore, even considering the much improved precision claimed in a recent study of LAGEOS II --- or, for that matter, any further refinement of $\|\alpha\| \ll 1$---, no qualitatively new results arise, since this only bridges the gap between $m $ and $M$ ({\it i.e.} narrows the value of $\|\xi - 1/12\|$). Finally, a word is due for the so-called chameleon mechanism, first posited in Ref. \cite{chameleon1,chameleon2,chameleon3,chameleon4}, as discussed in Ref. \cite{naf} in relation to $f(R)$ theories. This non-linear effect goes beyond the linear expansion of the modified field equations, and relies on the equivalence between $f(R)$ theories and a scalar-tensor theory with a scalar field $\phi$ proportional to $f_R$, which appears non-minimally coupled to the matter Lagrangian density (in the Einstein frame) \cite{analogy1,analogy2,analogy3,felice}. As it turns out, the effective potential of this scalar field can be written as $V_{eff}(\phi) = V(\phi) + e^{a\phi}\rho$ (with $a$ an appropriate constant), so that the position of its minimum depends on the density $\rho$, and the mass for the scalar field grows with the density: this is particularly relevant in a cosmological context, where the low background density yields a light, long-ranged field. Given the above, Ref. \cite{naf} speculates that further computations allowing for this non-linear effect could lead to different constraints on the mass scale $m$ of the adopted quadratic form for $f^1(R)$: quite naturally, the inclusion of a direct coupling between curvature ({\it vis-\`a-vis} the scalar field) and matter only heightens this possibility. Clearly, this prompts for a future study of the relation between this chameleon mechanism and a NMC model, in the framework of its equivalence with a multi-scalar-tensor theory \cite{multiscalar}.	14	3	1403.7251	In this work, the effects of a nonminimally coupled model of gravity on a perturbed Minkowski metric are presented. The action functional of the model involves two functions, f<SUP>1</SUP> (R) and f<SUP>2</SUP> (R), of the Ricci scalar curvature R: the former extends the usual linear term found in the Einstein-Hilbert Lagrangian, while the latter is multiplied by the matter Lagrangian density, thus introducing an explicit nonminimal coupling. <P />Based upon a Taylor expansion around R = 0 for both functions, we find that the metric around a spherical object is a perturbation of the weak-field Schwarzschild metric: the perturbation of the tt component of the metric tensor is shown to be a Newtonian plus Yukawa term, which can be constrained using the available experimental results. It is shown that this effect can be canceled or made arbitrarily small when the characteristic mass scales of the two functions are similar. We conclude that the Starobinsky model for inflation complemented with a generalized preheating mechanism is not experimentally constrained by observations. The geodetic precession effects of the model are also shown to be of no relevance for the constraints.	false	[ "R =", "an explicit nonminimal coupling", "the available experimental results", "the Ricci scalar curvature R", "Lagrangian", "observations", "Newtonian", "Yukawa", "the usual linear term", "the matter Lagrangian density", "a perturbed Minkowski metric", "Schwarzschild metric:", "the weak-field Schwarzschild metric", "Ricci", "f", "the characteristic mass scales", "Minkowski", "the metric tensor", "the tt component", "a Newtonian plus Yukawa term" ]	10.291885	0.588584	89
483048	[ "Lee, Chin-Fei", "Hirano, Naomi", "Zhang, Qizhou", "Shang, Hsien", "Ho, Paul T. P.", "Krasnopolsky, Ruben" ]	2014ApJ...786..114L	[ "ALMA Results of the Pseudodisk, Rotating Disk, and Jet in the Continuum and HCO<SUP>+</SUP> in the Protostellar System HH 212" ]	84	[ "-", "Academia Sinica Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 106, Taiwan", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Academia Sinica Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 106, Taiwan", "Academia Sinica Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 106, Taiwan; Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Academia Sinica Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 106, Taiwan" ]	[ "2014A&A...568L...5C", "2014ApJ...796...70C", "2015A&A...579A.114M", "2015A&A...581A..85P", "2015ASPC..499..205C", "2015ASPC..499..209P", "2015ASPC..499..219B", "2015Ap&SS.355..283B", "2015ApJ...799..193Y", "2015ApJ...802...86B", "2015ApJ...805..186L", "2015ApJ...812..129Y", "2015Natur.527...70P", "2016A&A...585A..74Y", "2016A&A...586L...3C", "2016A&A...590A..33P", "2016A&A...595L...4L", "2016A&ARv..24....6B", "2016ApJ...822...12H", "2016ApJ...826..213L", "2016ApJS..222....7L", "2016MNRAS.458.4222Z", "2016MNRAS.463.4246M", "2017A&A...606L...7B", "2017A&A...607L...6T", "2017ApJ...834..178Y", "2017ApJ...838...60M", "2017ApJ...839...47M", "2017ApJ...843...27L", "2017ApJ...843...45K", "2017ApJ...849...56A", "2017ApJ...850..158S", "2017NatAs...1E.152L", "2017SciA....3E2935L", "2018A&A...615A..58Y", "2018A&A...617A..10C", "2018ASSP...53..477S", "2018ApJ...856..164T", "2018ApJ...862....8T", "2018ApJ...864..168N", "2018ApJ...864L..25O", "2018IAUS..332..175Y", "2018MNRAS.473.3080M", "2018MNRAS.475..391M", "2018MNRAS.475.5322S", "2019A&A...629A..29J", "2019AJ....158..107R", "2019ApJ...871..100H", "2019ApJ...873...73Z", "2019ApJ...876...63L", "2019ApJ...876..149M", "2019ApJ...879..101L", "2019ApJ...887..209A", "2019ESC.....3.1564B", "2019ESC.....3.2110C", "2020A&A...635A..15M", "2020A&A...640A..82T", "2020ApJ...890..130T", "2020ApJ...891...61Y", "2020ApJ...893...54Y", "2020ApJ...905..116S", "2021A&A...650A.173S", "2021A&A...655A..65T", "2021ARep...65..693K", "2021ApJ...907L..41L", "2021ApJ...910...75L", "2021FrASS...8..137S", "2021MNRAS.501.1316L", "2022A&A...667A..20A", "2022ApJ...925...11D", "2022ApJ...925...12S", "2022ApJ...925...32T", "2022ApJ...937...10L", "2022MNRAS.515.6073K", "2022arXiv220913765T", "2023A&A...677A..92V", "2023ASPC..534..233P", "2023ASPC..534..317T", "2023ASPC..534..645P", "2023ApJ...946...70V", "2023INASR...8..144K", "2024ApJ...968...26T", "2024arXiv240315550T", "2024arXiv240508271M" ]	[ "astronomy" ]	15	[ "accretion", "accretion disks", "ISM: individual objects: HH 212", "ISM: jets and outflows", "stars: formation", "Astrophysics - Astrophysics of Galaxies" ]	[ "1976ApJ...210..377U", "1984ApJ...286..529T", "1990AJ.....99..924B", "1990ApJ...364..601R", "1992A&A...265..726Z", "1997A&A...318..595S", "1997ApJ...475..211O", "1998Natur.394..862Z", "1999ARA&A..37..311E", "2000ApJ...539..763G", "2000ApJ...543..291N", "2000ApJ...545.1034S", "2000prpl.conf..759K", "2000prpl.conf..789S", "2001ApJ...557..429L", "2002ApJ...580..987K", "2002MNRAS.336..797V", "2002Msngr.109...28M", "2003ApJ...599..363A", "2003MNRAS.339..633D", "2004MNRAS.351.1054R", "2006ApJ...639..292L", "2007A&A...462L..53C", "2007A&A...468L..29C", "2007ApJ...659..499L", "2008ApJ...681.1356M", "2008ApJ...685.1026L", "2009ApJ...699.1584L", "2009ApJ...700.1502A", "2010ApJ...723L..34C", "2010ApJ...725..712L", "2010MNRAS.406..102K", "2011ApJ...741...62L", "2012A&A...543A.128J", "2012A&A...548L...2C", "2012Natur.492...83T", "2013A&A...560A.103M", "2013ApJ...770..151C", "2013ApJ...774...82L", "2013MNRAS.435..289R", "2014A&A...562A..77H", "2014Natur.507...78S" ]	[ "10.1088/0004-637X/786/2/114", "10.48550/arXiv.1403.5853" ]	1403	1403.5853_arXiv.txt	Stars are formed inside molecular cloud cores by means of gravitational collapse. The details of the process, however, are complicated by the presence of magnetic fields and angular momentum. As a result, in addition to infall (or collapse), rotation and outflow are also seen toward star-forming regions. In theory, a rotationally supported disk (RSD) is expected to form inside a collapsing core around a protostar, from which part of the material is accreted by the protostar and part is ejected away. The RSD is expected to be Keplerian when the mass of the protostar dominates that of the disk. Observationally, however, when and how such a disk is actually formed are still unclear, because of the lack of detailed kinematic studies inside the collapsing core in the early phase of star formation. RSDs have been seen with a radius of $\sim$ 500 AU in the late (i.e., Class II or T Tauri) phase of star formation \cite[see, e.g.,][]{Simon2000}. Such disks must have formed early in the Class 0 phase, as claimed in a few Class 0 sources, e.g., HH 211 \cite[\mstar{0.05}, \rdisk{} $\lesssim$ 80 AU,][]{Lee2009}, NGC 1333 IRAS 4A2 \cite[\mstar{0.08}, \rdisk{} $\sim$ 130 AU,][]{Choi2010}, L 1527 \cite[\mstar{0.5}, \rdisk{} $\sim$ 125 AU,][]{Tobin2012}, and recently VLA 1623 \cite[\mstar{0.2}, \rdisk{} $\gtrsim$ 150 AU,][]{Murillo2013}. In models of non-magnetized core collapse, a RSD can indeed form as early as in the Class 0 phase \citep{Terebey1984}. However, a realistic model should include magnetic field, because recent survey toward a few Class 0 sources shows that molecular cores are magnetized and likely to have an hourglass B-field morphology \citep{Chapman2013}. Unfortunately, in many current models of magnetized core collapse, the magnetic field produces an efficient magnetic braking that removes the angular momentum and thus prevents a RSD from forming at the center \citep{Allen2003,Mellon2008}. In those cases, only a flattened envelope called the pseudodisk can be formed around the central source \cite[e.g.,][]{Allen2003}. Magnetic-field-rotation misalignment is sometimes able to solve this so-called magnetic braking catastrophe \citep{Joos2012,Li2013}, but not always. This paper is a follow-up study of the HH 212 protostellar system. This system is deeply embedded in a compact molecular cloud core in the L1630 cloud of Orion at a distance of 400 pc. The central source is the Class 0 protostar IRAS 05413-0104, with a bolometric luminosity $L_\textrm{\scriptsize bol}\sim$ 9 $L_\odot$ (updated for the distance of 400 pc) \citep{Zinnecker1992}. It drives a powerful bipolar jet \citep{Zinnecker1998,Lee2007}. A RSD must have formed in order to launch the jet, accordingly to current jet models. Previous observations in \cCO{} ($J=2-1$) and \bCO{} ($J=2-1$) with the Submillimeter Array (SMA) at $\sim$ \arcsa{2}{5} resolutions showed a flattened envelope around the central source \citep{Lee2006}. It is infalling with small rotation (i.e., spiraling) into the central source, and thus can be identified as a pseudodisk. The derived infall velocity and rotation velocity in the flattened envelope suggested a small RSD to be present at the center with a radius of $\sim$ 100 AU. However, previous \HCOP{} J=4-3 observations toward the inner part of the flattened envelope with the SMA at $\sim$ \arcs{1} resolution failed to confirm the presence of such a disk \citep{Lee2007}. Later observations at higher angular resolution in continuum suggested again a compact disk around the central source \citep{Codella2007,Lee2008}. In order to confirm the presence of a RSD in this system, we have mapped this system in 350 GHz continuum and \HCOP{} J=4-3 using Atacama Large Millimeter/Submillimeter Array (ALMA) at higher resolution and sensitivity. As before, we can identify a disk around the central source inside the flattened envelope. We model the continuum and \HCOP{} emission simultaneously, deriving the kinematic and physical properties of both the flattened envelope and the disk. The disk indeed could have a Keplerian rotation profile. In addition, in our model, in order to produce enough continuum emission of the disk, a density jump is required across the interface between the envelope and the disk. We will discuss how the disk can be formed inside the envelope. Since \HCOP{} also traces the jet close to the central source, we will also discuss the properties and origin of the \HCOP{} emission in the jet.	We have mapped the HH 212 protostellar system in the dust continuum at 350 GHz and the \HCOP{} (J=4-3) line. Our primary conclusions are the following: \begin{itemize} \item A flattened envelope is seen in continuum and \HCOP{} around the central source, extending out to $\sim$ 800 AU (\arcs{2}). The \HCOP{} kinematics shows that the flattened envelope is infalling with rotation (i.e., spiraling) into the central source, and thus can be identified as a pseudodisk found in the models of magnetized core collapse. \item Inside the flattened envelope, a bright compact disk is seen in continuum at the center with an outer radius of $\sim$ 120 AU (\arcsa{0}{3}). This disk is also seen in \HCOP{} and the \HCOP{} kinematics shows that the disk is rotating. \item \HCOP{} emission is missing at low-redshifted velocity centered at $\Voff \sim$ 0.2 \vkm{}. To account for this missing, an extended infalling envelope is required, with its material flowing roughly parallel to the jet axis toward the pseudodisk. This is expected if the envelope is magnetized with an hourglass B-field morphology. Since the material can flow along the field lines and the field lines are roughly parallel to the jet axis, the material can flow parallel to the jet axis into the pseudodisk. \item We have modeled the continuum and \HCOP{} emission of the flattened envelope and disk simultaneously. In our model, we assume a change of rotation profile from that of conserving specific angular momentum in the flattened envelope to that of Keplerian in the disk, as inspired by theoretical models of core collapse and previous observations of HH 111. In order to reproduce enough disk continuum emission at the center, a jump in density with a factor of $\sim$ 8 is required across the interface between the flattened envelope and the disk. This jump turns out to be consistent with an isothermal shock at the interface. In our model, the disk only has $\sim$ 8\% of the stellar mass and thus can indeed be formed. However, further observations at higher resolution are needed to confirm its Keplerian rotation and its radius. \item A collimated jet is seen in \HCOP{} extending out to $\sim$ 500 AU from the central source (and disk?), with its peaks located upstream of those seen before in SiO. The \HCOP{} emission is seen with a broad range of velocities. The \HCOP{} abundance is highly enhanced by a factor of $\sim$ 100, comparing to that of the flattened envelope. All these suggest that the \HCOP{} emission traces internal shocks in the jet. \end{itemize}	14	3	1403.5853	HH 212 is a nearby (400 pc) Class 0 protostellar system showing several components that can be compared with theoretical models of core collapse. We have mapped it in the 350 GHz continuum and HCO<SUP>+</SUP> J = 4-3 emission with ALMA at up to ~0.''4 resolution. A flattened envelope and a compact disk are seen in the continuum around the central source, as seen before. The HCO<SUP>+</SUP> kinematics shows that the flattened envelope is infalling with small rotation (i.e., spiraling) into the central source, and thus can be identified as a pseudodisk in the models of magnetized core collapse. Also, the HCO<SUP>+</SUP> kinematics shows that the disk is rotating and can be rotationally supported. In addition, to account for the missing HCO<SUP>+</SUP> emission at low-redshifted velocity, an extended infalling envelope is required, with its material flowing roughly parallel to the jet axis toward the pseudodisk. This is expected if it is magnetized with an hourglass B-field morphology. We have modeled the continuum and HCO<SUP>+</SUP> emission of the flattened envelope and disk simultaneously. We find that a jump in density is required across the interface between the pseudodisk and the disk. A jet is seen in HCO<SUP>+</SUP> extending out to ~500 AU away from the central source, with the peaks upstream of those seen before in SiO. The broad velocity range and high HCO<SUP>+</SUP> abundance indicate that the HCO<SUP>+</SUP> emission traces internal shocks in the jet.	false	[ "magnetized core collapse", "core collapse", "high HCO", "HCO", "<", "theoretical models", "disk", "emission", "internal shocks", "~500 AU", "the central source", "small rotation", "SiO.", "resolution", "several components", "an extended infalling envelope", "the missing HCO<SUP>+</SUP> emission", "the jet axis", "a compact disk", "ALMA" ]	10.62307	11.529901	-1
769602	[ "Vereshchagin, S. V.", "Chupina, N. V.", "Sariya, Devesh P.", "Yadav, R. K. S.", "Kumar, Brijesh" ]	2014NewA...31...43V	[ "Apex determination and detection of stellar clumps in the open cluster M 67" ]	15	[ "Institute of Astronomy, Russian Academy of Sciences (INASAN), 48 Pyatnitskaya st., Moscow, Russia", "Institute of Astronomy, Russian Academy of Sciences (INASAN), 48 Pyatnitskaya st., Moscow, Russia", "Aryabhatta Research Institute of Observational Sciences, Manora Peak, Nainital 263 002, India; School of Studies in Physics & Astrophysics, Pt. Ravishankar Shukla University, Raipur 492 010, CG, India", "Aryabhatta Research Institute of Observational Sciences, Manora Peak, Nainital 263 002, India;", "Aryabhatta Research Institute of Observational Sciences, Manora Peak, Nainital 263 002, India;" ]	[ "2015Ap.....58..522E", "2015MNRAS.452.3394M", "2016A&A...592L...1B", "2016BaltA..25..432V", "2016NewA...49...32E", "2017Ap.....60..173H", "2018Ap&SS.363...58E", "2020AJ....160..119B", "2020ARep...64...94E", "2020RAA....20...16P", "2021AJ....161..101S", "2021AJ....162..146S", "2021PASJ...73..677B", "2022AJ....164..171B", "2024AJ....167..188B" ]	[ "astronomy" ]	7	[ "Astrophysics - Astrophysics of Galaxies" ]	[ "1969AJ.....74....2V", "1977A&AS...27...89S", "1984MNRAS.206..529E", "1986AJ.....92.1100M", "1989AJ.....98..227G", "1991A&A...244...69B", "1993AJ....106..181M", "1997A&A...323L..49P", "1998A&A...334..552C", "2000SPIE.4008..467B", "2001A&A...371..115C", "2001A&A...375..851M", "2004MNRAS.347..101S", "2005A&A...438.1163K", "2005A&A...439..805V", "2006A&A...450..557R", "2006A&A...454.1029A", "2008A&A...484..609Y", "2008MNRAS.390..665L", "2008MNRAS.391..343P", "2009A&A...503..165Y", "2009ApJ...698.1872S", "2010A&A...513A..51B", "2011JAVSO..39..219M", "2012A&A...543A..87S", "2012A&A...545A.139P", "2013MNRAS.430.3350Y" ]	[ "10.1016/j.newast.2014.02.008", "10.48550/arXiv.1403.2532" ]	1403	1403.2532_arXiv.txt	\label{Intro} M~67 is one of the most studied open clusters among the known open clusters with ages comparable to or older than the Sun. M~67 is of Solar metallicity, relatively nearby and has low interstellar reddening. It has been comprehensively studied by many authors to establish astrometric membership \citep{Sanders1977, Girard1989,Yadav2008}. Many photometric studies \citep{Montgomery1993, Sandquist2004} and rather precise radial velocity and binary search study \citep{Mathieu1986, Melo2001, Pasquini2012} have been conducted for the cluster. The fundamental parameters along with absolute proper motions have been listed in Table 1. M~67 moves far above the galactic disk on latitude $b=+31.91^\circ$, on $z=830\sin(31.91^\circ)=440$~pc along quite a circular orbit and interacts with spiral density waves, which can initiate star formation. Using 2MASS $JHK$ photometry, \citet{Sarajedini2009} suggested an age of 3.5 Gyr using two different theoretical isochrones. \begin{table} \caption{The fundamental parameters for M~67} \vspace{0.5cm} \centering \scriptsize \begin{tabular}{lll} \hline\hline Parameters&Value &Ref.\\ \hline \noalign{\smallskip} Equatorial coordinate (J2000)&$\alpha=8.855^h$, $\delta=11.8^\circ$ & WEBDA\\ Galactic coordinate &$l=215.69^\circ$, $b=+31.91^\circ$& WEBDA\\ The core radius &$r_c=0.12^\circ$&Kharchenko et al. (2005)\\ The cluster radius &$r_{cl}=0.78^\circ$ &Kharchenko et al. (2005)\\ Distance from the Sun &830 pc& \citet{Allen1973} \\ Age &3.5--4.0 Gyr&\citet{Sarajedini2009} \\ $[$Fe/H$]$ &+0.03$\pm$0.01&\citet{Randich2006}\\ Absolute proper motion &$\mu_\alpha\cos\delta=-9.6\pm1.1,\, \mu_\delta=-3.7\pm0.8$ mas/yr &\citet{Bellini2010} \\ \hline \label{tab1} \end{tabular} \end{table} \citet{Chupina1998} have detected several stellar clumps inside the low density extended corona of M~67 with the help of nearest neighbour distance (NND) method, while the cause of origin of these substructures remains unclear. Availability of new proper motions and radial velocities data for M~67 prompted us to revise the membership of stars in the cluster with our methods using radial velocities and proper motion data to ascertain membership. The stellar apexes diagram or $AD$-diagram, \citep{Chupina2001} is useful for the investigation of kinematic structures of the star clusters and streams. It allows us to find the kinematic substructures inside these objects. The AD-diagram is the plot of the individual star apexes. The individual apexes represent the equatorial coordinates of the point on the celestial sphere in which the space velocity vector intersects it. By analogy, the star apex coordinates of the cluster are designated in equatorial coordinates as $A$ for right ascension and $D$ for declination. The formal description of this method and the formulae of the error ellipses are given in \citet{Chupina2001}. It should be noted that the error ellipses can be constructed only using the Hipparcos data because it contains the necessary correlation coefficients. The main purpose of the present analysis is the membership revision of \citet{Yadav2008} catalogue with the help of convergent point method. For this purpose, we determined apex coordinates and used the methods from our previous work \citep{Chupina2001}. The new data has also been used to study the substructures in its central part. This work complements our previous one on the corona of the cluster \citep{Chupina1998}. \begin{table} \caption{Characteristics of the data taken from \citet{Yadav2008} catalogue.} \vspace{0.5cm} \centering \begin{tabular}{ccccc} \hline\hline V range& $\sigma_{\mu_{\alpha}}$& $\sigma_{\mu_{\delta}}$ & $\sigma_{V_r}$ & N\\ \hline \noalign{\smallskip} 7-10 & 1.86 & 3.27 & 0.008 & 3 \\ 10 - 13 & 2.48 & 2.61 & 0.099 & 120 \\ 13 - 16 & 2.24 & 2.99 & 0.156 & 451 \\ 16 - 19 & 3.60 & 4.13 & 0.116 & 570 \\ 19 - 22 &15.16 &15.52 & -- & 1266 \\ \hline \label{tab2} \end{tabular} \end{table} The structure of the article is as follows: The data used for the present analysis is described in Sect.~2, while Sect.~3 is devoted to apex determination. Comments on membership are presented in Sect.~4, while Sect.~5 describes the substructures in the central part of the cluster. Finally, in Sect. 6 we list the conclusions of the present analysis.	From the present analysis for M~67 open cluster, the following main conclusions can be drawn: \begin{enumerate} \item The apex coordinates for M~67 have been calculated as: $A_0=132.97^\circ\pm0.81^\circ,\, D_0=11.85^\circ\pm0.90^\circ$. \item The membership of the stars has been revisited. In future, it is possible to decide membership by distance from $\mu$-box and by position of the star in ``$\mu_U-\mu_T$''-diagram. If $\mu_T$ is close to zero, the membership probability is higher (within errors and peculiar velocities dispersion). \item AD-diagram has been constructed for stars with known radial velocities. The membership status for the star can be estimated with the help of AD-diagram. The farther away a star located from the apex position, there will be the less probability for it to be a cluster member. \item We found substructures in the corona of M~67. We found heterogeneity in the core periphery, like our previous work \citep{Chupina1998}. \end{enumerate}	14	3	1403.2532	We determined the cluster’s apex coordinates, studied the substructures and performed membership analysis in the central part (34<SUP>‧</SUP>×33<SUP>‧</SUP>) of the open cluster M 67. We used the individual stellar apexes method developed earlier and classical technique of proper motion diagrams in coordinate system connected with apex. The neighbour-to-neighbour distance technique was applied to detect space details. The membership list was corrected and some stars were excluded from the most probable members list. The apex coordinates have been determined as: A<SUB>0</SUB>=132.97° ± 0.81° and D<SUB>0</SUB>=11.85° ± 0.90°. The 2D-space star density field was analysed and high degree of inhomogeneity was found.	false	[ "° ±", "±", "coordinate system", "apex", "proper motion diagrams", "high degree", "membership analysis", "space details", "D<SUB>0</SUB>=11.85", "inhomogeneity", "The apex coordinates", "the individual stellar apexes method", "The membership list", "the cluster’s apex coordinates", "the most probable members list", "The 2D-space star density field", "cluster M 67", "the central part", "distance", "some stars" ]	8.209005	9.449058	-1
746247	[ "Munshi, D.", "Corasaniti, P. S.", "Coles, P.", "Heavens, A.", "Pandolfi, S." ]	2014MNRAS.442.3427M	[ "Reionization and CMB non-Gaussianity" ]	3	[ "School of Mathematical and Physical Sciences, University of Sussex, Brighton BN1 9QH, UK; School of Physics and Astronomy, Cardiff University, Queen's Buildings, 5 The Parade, Cardiff CF24 3AA, UK", "Laboratoire Univers et Théories (LUTh), UMR 8102 CNRS, Observatoire de Paris, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France", "School of Mathematical and Physical Sciences, University of Sussex, Brighton BN1 9QH, UK; School of Physics and Astronomy, Cardiff University, Queen's Buildings, 5 The Parade, Cardiff CF24 3AA, UK", "Imperial Centre for Inference and Cosmology, Blackett Laboratory, Imperial College, Prince Consort Road, London SW7 2AZ, UK", "Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries, Vej 30, 2100 Copenhagen, Denmark" ]	[ "2017JCAP...02..010M", "2017PhRvD..95h3512C", "2018MNRAS.481..970R" ]	[ "astronomy" ]	3	[ "methods: analytical", "methods: numerical", "methods: statistical", "dark ages", "reionization", "first stars", "diffuse radiation", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1998ApJ...496..605E", "1998MNRAS.299..805H", "1999ApJ...527L...1W", "2000ApJ...529...12H", "2001ARA&A..39...19L", "2001ApJ...561..504L", "2001PhR...340..291B", "2001PhRvD..64d3516C", "2003ApJ...583...24K", "2003ApJ...595...13H", "2003ApJ...598..756S", "2003PhRvD..68b3001H", "2004ApJ...600...26H", "2004PhR...402..103B", "2004PhRvD..69b3512K", "2004PhRvD..70b3508C", "2004PhRvD..70d3502C", "2005ApJ...634...14K", "2006AJ....132..117F", "2006ARA&A..44..415F", "2006ApJ...647..840A", "2006JCAP...05..004C", "2006MNRAS.373..561L", "2007ApJ...663L...1H", "2007JCAP...03..019C", "2007PhRvD..76d3002D", "2007PhRvD..76d3510S", "2008ApJ...672..737M", "2008MNRAS.384..291A", "2008PhRvL.100r1301Y", "2009PhRvD..79b3501K", "2009PhRvD..79d3003D", "2009PhRvD..79j7302D", "2010MNRAS.401.2406M", "2010MNRAS.402.2617T", "2010PhRvD..81l3509P", "2010PhRvD..82l3527P", "2011MNRAS.414.3173M", "2011MNRAS.417....2S", "2011PhRvD..84l3522P", "2012MNRAS.419..138M", "2012MNRAS.419..536M", "2013ApJ...769...93P", "2013ApJ...779..124M", "2013ApJS..208...20B", "2013MNRAS.428.2628M", "2013MNRAS.434.2830M", "2014A&A...571A..16P", "2014A&A...571A..17P", "2014A&A...571A..22P", "2014A&A...571A..24P" ]	[ "10.1093/mnras/stu1123", "10.48550/arXiv.1403.1531" ]	1403	1403.1531_arXiv.txt	Pinning down the details and controlling physics of the cosmological reionization history remains one of the important goals of present-day cosmology. It is well known, thanks to a large set of astrophysical observables, that after primordial recombination, which occurred at a redshift of $z \sim 1100$, the Universe reionized at a redshift $z>6$. The epoch-of-reionization (EoR) is related to many fundamental questions in cosmology, such as the properties of the first galaxies, physics of (mini-)quasars, formation of very metal-poor stars and a slew of other important research topics in astrophysics. Hence uncovering it will have far reaching implications on the study of structure formation in the early Universe \citep{LB01}. Observations of Lyman-$\alpha$ forests with high-resolution echelle spectrographs on large telescopes (such as HIRES on Keck, and UVES on ESO's Very Large Telescope) are valuable for studying reionization at $z\approx 2.5-6.5$ \citep{Fan06,FCK06}. Redshifted 21cm observations are also a very important probe of the EoR and several instruments are either operational or in the construction phase. In the short term these consists of: The Low Frequency Array (LOFAR)\footnote{http://www.lofar.org/}, the Murchison Widefield Array (MWA)\footnote{http://www.mwatelescope.org/}, Precision Array to Probe Epoch of Reionization (PAPER) and Giant Metrewave Radio Telescope (GMRT)\footnote{http://gmrt.ncra.tifr.res.in/}, while, on a somewhat longer time scale the Square Kilometre Array (SKA)\footnote{http://www.skatelescope.org/} will be operational. In addition to Lyman-$\alpha$ and 21cm redshifted observations, CMB temperature and polarisation studies can also provide valuable information regarding the EoR. The polarisation signal in CMB is generated due to the scattering of the local CMB temperature quadrupole by the free-electron population. This signal peaks at angular scales corresponding to the horizon at the rescattering surface (at a few tens of degrees) and the amplitude depends on total optical depth. However, the large cosmic variance associated with the signal means it is impossible to discriminate among various reionization histories using cross-correlation of the CMB temperature and polarisation \citep{Kap03,HZKK08,HuHo08}. The total optical depth to reionization using WMAP data is $\tau=0.08 \pm 0.013$ \citep{Bennett13}. Most current constraints from CMB data are analysed assuming a ``sudden'' and complete reionization at a redshift $z_r$ for WMAP\footnote{http://map.gsfc.nasa.gov/} this value of $\tau$ will correspond to $z_r=11$. However as mentioned before, the precise details of the reionization process are not very well known and clearly the reionization history of the universe at those redshifts could have easily been very different. The combination of temperature data and lensing reconstruction from the Planck data gives an optical depth $\tau=0.089 \pm 0.032$ \citep{Planck13d}\footnote{http://www.rssd.esa.int/index.php?project=Planck}, consistent with WMAP9 estimates. The polarization data from Planck is expected to improve the accuracy of determination of $\tau$, however, it is important to keep in mind that CMB observations only provide integrated or projected information on reionization. The process of reionization is expected to be patchy and inhomogeneous in scenarios where reionization is caused by UV emission from the first luminous objects \citep[e.g.][]{Meerburg13}, and the resulting fluctuations in visibility will generate extra anisotropy at arcminute scales. Even in scenarios when reionization is caused by energy injection from decaying particles or X-ray emission, inhomogeneities in electron density can cause fluctuations in visibility, but these fluctuations are too small to be detected in temperature and polarization power spectra or their cross-spectrum. Nevertheless, as one can imagine, additional signals due to inhomogeneneites in the free-electron density may imprint additional features at smaller angular scales \citep{Hu00,Santos}, and the mixed bispectrum with external data sets was advocated to extract redshift information \citep{Cooray04, Al06,AF07,T10,HIM07}. In fact, what is required is the three-point correlation (or equivalently the bispectrum) involving temperature, polarisation and a tracer field for the free electron population can extract useful information on the ionization history of the Universe \citep{Cooray04}. For details of the generation of secondary non-Gaussianity due to reionization see \citep{KW10}. The estimation of the bispectrum is a lot more complicated than the power spectrum due to the presence of additional degrees of freedom. As has been pointed out in many recent works, the mode-by-mode estimation of the bispectrum, though very attractive, is seldom useful because of the associated low signal-to-noise \citep[e.g.][for a review]{Bart04}. Typically, this means, the entire information content of the bispectrum is often compressed into a single number which is used to distinguish various models of reionization. Though this has the advantage of increasing the (S/N), it also degrades the information content of the bispectrum. A compromise solution was proposed recently by \cite{MuHe10}, who defined a power spectrum associated with a specific bispectrum. This power spectrum represents the cross-spectra of the product of two maps $[\rm X(\oh)Y(\oh)]$ against another map $\rm Z(\oh)$. It is a weighted sum of individual modes of the bispectrum keeping one of the index fixed while summing over the other two indices. Such an estimator can also be designed to work with an experiential mask for estimation in the presence of non-uniform noise and is relatively simple to implement. In the literature such estimators are known as pseudo-$\myC_\ell$ (PCL) estimators. However, such estimators are sub-optimal. With the recent attempts to detect primordial non-Gaussianity in the aftermath and leading up to the Planck data release \citep{Planck13a,Planck13b} there has been an increased activity in the area of optimising estimators which can probe primordial non-Gaussianity \citep{Heav98,KSW,Crem06,Crem07b,SmZaDo00,SmZa06}. Detection of secondary non-Gaussianity, can also benefit from using the \cite{MuHe10} estimators to take into account inhomogeneous noise and partial sky coverage in an optimal way \citep{MuCoHeVa11}. We will discuss these issues and other related optimisation problems in this paper for mixed data sets. Being able to probe the bispectrum in a scale-dependent way will provide a useful way to differentiate among different theories of reionization. The matched filtering inherent in these estimators are likely to be very useful in pinning down a specific reionization history. The formalism also provides a natural set-up to study cross-contamination from effects of weak lensing. The approach presented here has already been successfully implemented for Planck data analysis which resulted in the detection of non-Gaussianity from the correlation of the Integrated Sachs-Wolfe (ISW) effect with gravitational lensing, and from residual point sources in the maps \citep{Planck13b}. We will define several set of different estimators. In addition to using the direct estimators, we will define, three dimensional constructs that require the use of appropriate weight functions to cross-correlate and probe the secondary non-Gaussianity. We will also point to computationally extensive estimators which can take into account all possible complications. These generalisations involve a set of fully optimal estimators which can work directly with harmonics of associated fields and carry out inverse covariance weighting using a direct brute force approach. Clearly, though such a direct approach is completely optimal it is prohibitively expensive to implement beyond a certain resolution. Nevertheless, for secondary bispectrum which we consider here, it will be important to maintain the optimality to a high resolution as most relevant information will be appear on small scales. The results presented here can be seen as an extension of our earlier papers: for example, \citep{MuHe10}, where we presented skew-spectrum for primary non-Gaussianity; \citep{MuCoHeVa11}, where results relevant to secondary non-Gaussianity were obtained. Here, we include polarization data in addition to the temperature maps and cross-correlate with external data sets in 3D to constrain various scenarios of reionization. In recent years we have extended the concept of skew-spectra to study topological properties of CMB maps \citep{MuCoHe13, MuSmCoReHaCo13} as well as for other cosmological data sets such as frequency-cleaned thermal Sunyaev Zeldovich $y$-maps \citep{MuSmJoCo12a} or weak lensing maps \citep{MuSmWaCo12b}. This paper is organised as follows: In \textsection\ref{ana} we present many of our analytical results and include the description of reionization models and the tracer fields that we study. We introduce our estimators in \textsection\ref{sec:estim}. In \textsection\ref{sec:disc} we discuss our results and \textsection\ref{sec:conclu} is devoted to concluding remarks. In Appendix \ref{sec:patch} we outline how a skew-spectrum estimator can be constructed using minimum variance estimated of fluctuations in optical depth. In Appendix \ref{sec:lensing_recon} we provide equivalent estimators for the reconstruction of the lensing potential.	\label{sec:conclu} The free electron population during reionization epoch re-scatters the local CMB temperature quadrupole and generates an additional polarization signal at small angular (arcminutes) scale. Due to their small amplitude, this contribution cannot be studied using the CMB temperature-polarization cross-spectra. However, additional information regarding the temporal evolution of the spatial variation of the free electron density can be gained by studying the three-way correlation between temperature anisotropy, polarisation and an external field, which can act as a tracer field for the free electron density. In harmonic space the associated mixed bispectrum can be used to constrain models of reionization. Estimation of individual modes of bispectra are dominated by noise, so the majority of studies in the past have used the skewness, which compresses all available modes to a single high (S/N) number, but this may mask the reionization history. Here we have shown how the recently proposed skew-spectra can be used to discriminate between models of reionization. We find that the amplitude of the skew-spectra correlates strongly with the epoch of reionization as well as the redshift distribution of tracers. We have studied four different models of reionization and three realistic tracers. We find that the use of multiple tracers can be very powerful in probing the redshift evolution of the ionization fraction. Most of the signal comes from high $\ell$ hence surveys and tracers with limited sky coverage can also provide valuable information and all-sky coverage is not an absolute necessity. Our results correspond to a Planck-type beam but experiments with even higher angular resolution will be able to achieve higher (S/N). In principle with judicious choice of different tracers it will be possible to map out the entire ionization history. We develop both the direct or PCL-based estimators as well as inverse covariance-weighted optimal estimators. For each choice of tracer field, we develop three set of estimators for cross-validation in Eq. (\ref{eq:diff_skew}). \cb{The contamination from weak-lensing was found to be negligible. In case of an ideal experiment without detector noise, depending on the redshift distribution of the tracer field, the (S/N) for detection of skew-spectra can reach relatively high values in most scenarios typically ${\cal O}(20)$, and even higher for some scenarios. For Planck type experiment the (S/N) for most scenarios is typically ${\cal O}(10)$.} In the text of the paper we have used the visibility function as our primary variable to describe the reionization history. In the Appendix we detail equivalent results for optical depth instead .The patchy reionization induces non-Gaussianity both due to patchy screening as well as Thomson scattering. We show that the estimators for reconstructing fluctuations in optical depth can be cross-correlated with external data sets, and the resulting estimators are similar to the skew-spectra but with different weights. These estimators that work with minimum variance reconstruction of optical depth are however are not optimal and differ from the corresponding PCL estimators. Cross-correlating with tracers which have redshift information has the advantage of distinguishing different histories of reionization. It is generally believed that the polarization from late-time reionization by patchy screening is small compared with polarization due to late-time Thomson scattering during reionization. However recent studies based on power-spectrum analysis have shown that at small angular scales both effects are comparable \citep{DHS09}. We derive the bispectrum generated by patchy screening of primary as well as by late-time Thomson scattering. The primary motivation of this paper was to devise a method to distinguish between different reionization histories of the Universe, using non-Gaussianity induced by the fluctuations in optical depth. This has been achieved, using mixed bispectra of CMB temperature and polarisation fields, along with one or more foreground tracers of free electron density. Using skew-spectra, originally devised for studies of primordial non-Gaussianity, we find that different reionization history models can be distinguished with this method with high signal-to-noise. \cb{We have assumed a perfect subtraction of all foregrounds to arrive at our results. Needless to say, that, as in any study using CMB data, unsubtracted residuals from the component separation step of the data reduction pipeline can seriously bias conclusion drawn using techniques presented here.}	14	3	1403.1531	We show how cross-correlating a high-redshift external tracer field, such as the 21-cm neutral hydrogen distribution and product maps involving cosmic microwave background (CMB) temperature and polarization fields, that probe mixed bispectrum involving these fields, can help to determine the reionization history of the Universe, beyond what can be achieved from cross-spectrum analysis. Taking clues from recent studies for the detection of primordial non-Gaussianity, we develop a set of estimators that can study reionization using a power spectrum associated with the bispectrum (or skew-spectrum). We use the matched filtering inherent in this method to investigate different reionization histories. We check to what extent they can be used to rule out various models of reionization and study cross-contamination from different sources such as the lensing of the CMB. The estimators can be fine-tuned to optimize study of a specific reionization history. We consider three different types of tracers in our study, namely: proto-galaxies; 21-cm maps of neutral hydrogen; and quasars. We also consider four alternative models of reionization. We find that the cumulative signal-to-noise ratio (S/N) for detection at ℓ<SUB>max</SUB> = 2000 can reach O(70) for cosmic variance limited all-sky experiments. Combining 100 GHz, 143 GHz and 217 GHz channels of the Planck experiment, we find that the S/N lies in the range O(5)-O(35). The S/N depends on the specific choice of a tracer field, and multiple tracers can be effectively used to map out the entire reionization history with reasonable S/N. Contamination from weak lensing is investigated and found to be negligible, and the effects of Thomson scattering from patchy reionization are also considered.	false	[ "different reionization histories", "patchy reionization", "reionization", "cross-spectrum analysis", "multiple tracers", "neutral hydrogen", "cosmic microwave background", "tracers", "product maps", "a specific reionization history", "the entire reionization history", "cosmic variance", "CMB", "different sources", "the reionization history", "recent studies", "study", "a tracer field", "reasonable S/N. Contamination", "CMB) temperature and polarization fields" ]	13.127893	3.187075	-1
696950	[ "Plazas, A. A.", "Bernstein, G. M.", "Sheldon, E. S." ]	2014JInst...9C4001P	[ "Transverse electric fields' effects in the Dark Energy Camera CCDs" ]	17	[ "Department of Physics, Brookhaven National Laboratory, Bldg. 510, Upton NY, 11973, USA", "Department of Physics and Astronomy, University of Pennsylvania, Davind Rittenhouse Laboratory, 209 South St., Philadelphia PA, 19104, USA", "Department of Physics, Brookhaven National Laboratory, Bldg. 510, Upton NY, 11973, USA" ]	[ "2014SPIE.9150E..17R", "2015ApJ...805...40M", "2015JInst..10C5013A", "2015JInst..10C5015W", "2015JInst..10C5017M", "2015JInst..10C5027B", "2015JInst..10C6010N", "2018MNRAS.481.1149Z", "2019ApJ...874..106B", "2019PhRvL.122p1801A", "2021MNRAS.501.1282J", "2021MNRAS.504.4312G", "2022AJ....164..128Z", "2023PASP..135k5003E", "2023SCPMA..6629811X", "2024arXiv240114944L", "2024arXiv240606472L" ]	[ "astronomy", "physics" ]	11	[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]	[ "1995A&AS..113..587M", "1996A&AS..117..393B", "2005astro.ph.10346T", "2006ASPC..351..112B", "2006SPIE.6269E..3KE", "2008arXiv0810.3600H" ]	[ "10.1088/1748-0221/9/04/C04001", "10.48550/arXiv.1403.6127" ]	1403	1403.6127_arXiv.txt	\label{sec:intro} In the past decade, the development of thick, high-resistivity CCDs with high quantum efficiency (QE) at long wavelengths (near infrared) has been encouraged by increasing scientific interest in this part of the electromagnetic spectrum {\citep{holland2003,holland2007}}. Thick CCDs also offer other advantages over more conventional and thin CCDs by reducing fringing at long wavelengths. Thus, thick, fully-depleted CCDs have been chosen by several current and future astronomical surveys {and imagers} (\emph{e.g.}, the \emph{Dark Energy Survey}, DES\footnote{\url{www.darkenergysurvey.org}} {\citep{abbott2005}}, {the \emph{Pan-STARRS} survey \citep{kaiser2010}, the \emph{Hyper Suprime Camera} \citep{komiyama2010}}, and the \emph{Large Synoptic Survey Telescope}, LSST\footnote{\url{www.lsst.org}} {\citep{ivezic2008}}). {Despite their advantages, thick CCDs also exhibit undesirable characteristics that can leave signatures in the data, with consequences for shape, point spread function (PSF), astrometric, and photometric measurements \citep{stubbs2014,lupton2014,antilogus2014,jarvis2014}}. {Current and future experiments, such as DES and LSST, will produce large data sets, which will provide higher statistical precision. Thus, it has become} necessary to fully characterize and understand even the most subtle systematic effects in the methods and instrumentation {\citep{weinberg2013}}. For instance, supernovae probes are limited by color and flux calibration, not statistical errors. In addition, weak gravitational lensing (WL) of large scale structure of the Universe (cosmic shear) requires measuring galaxy shapes and positions with high accuracy in order to exploit its full potential to constrain dark energy \citep{albrecht2006}. In this paper, we study the consequences of electric fields transverse to the surface of the CCDs of the Dark Energy Camera (DECam, {\cite{diehl2012,flaugher2012}; Flaugher et al. (in prep.)}), the imager built for the DES project. These transverse fields\footnote{Also referred to as ``lateral fields".} have an impact on astrometric and photometric measurements by shifting the location of the collected charge and modifying the effective pixel area. Charge relocation and pixel area variations contribute to the local variations in the response to uniform illumination (Pixel Response Non-Uniformity, PRNU), which are usually assumed to originate purely from sensitivity (or QE) differences in the pixels. Thus, naively dividing by the flat-field images could lead to systematic errors in the calibration steps during the data reduction process. In the next Section, we provide a brief introduction to DECam in the context of the DES and illustrate some of the structures visible in flat-field images that do not correspond purely to pixel sensitivity differences. In Section $\S$3, we use flat-field images to derive templates of the amplitude of these effects as a function of their pixel location in the detectors. We then demonstrate, in Section $\S$4, how these templates can be incorporated into astrometric and photometric solutions to reduce residuals below the scientific requirements for DES. We conclude and summarize our results and their implications in Section $\S$5.		14	3	1403.6127	Spurious electric fields transverse to the surface of thick CCDs displace the photo-generated charges, effectively modifying the pixel area and producing noticeable signals in astrometric and photometric measurements. We use data from the science verification period of the Dark Energy Survey (DES) to characterize these effects in the Dark Energy Camera (DECam) CCDs, where the transverse fields manifest as concentric rings (impurity gradients or ``tree rings'') and bright stripes near the boundaries of the detectors (``edge distortions'') with relative amplitudes of about 1% and 10%, respectively. Using flat-field images, we derive templates in the five DES photometric bands (grizY) for the tree rings and the edge distortions as a function of their position on each DECam detector. Comparison of the astrometric and photometric residuals confirms their nature as pixel-size variations. The templates are directly incorporated into the derivation of photometric and astrometric residuals. <P />The results presented in these proceedings are a partial report of analysis performed before the workshop ``Precision Astronomy with Fully depleted CDDs'' at Brookhaven National Laboratory. Additional work is underway, and the final results and analysis will be published elsewhere (Plazas, Bernstein & Sheldon 2014, in prep.).	false	[ "edge distortions", "10%", "Brookhaven National Laboratory", "concentric rings", "relative amplitudes", "noticeable signals", "about 1% and 10%", "DECam", "about 1%", "impurity gradients", "bright stripes", "prep", "Spurious electric fields", "pixel-size variations", "photometric and astrometric residuals", "thick CCDs", "astrometric and photometric measurements", "DES", "the edge distortions", "analysis" ]	10.82258	3.892094	171
437260	[ "Vaduvescu, O.", "Kehrig, C.", "Bassino, L. P.", "Smith Castelli, A. V.", "Calderón, J. P." ]	2014A&A...563A.118V	[ "Searching for star-forming dwarf galaxies in the Antlia cluster" ]	8	[ "Isaac Newton Group of Telescopes, Apto. 321, 38700, Santa Cruz de la Palma, Canary Islands, Spain", "Instituto de Astrofísica de Andalucía (CSIC), Apto. 3004, 18080, Granada, Spain", "Grupo de Investigación CGGE, Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo del Bosque, B1900FWA, La Plata, Argentina; Consejo Nacional de Investigaciones Científicas y Técnicas (CONICET), C1033, AAJ Ciudad Autónoma de Buenos Aires, Argentina; Instituto de Astrofísica de La Plata (CCT-La Plata, CONICET-UNLP), Paseo del Bosque, B1900FWA, La Plata, Argentina", "Grupo de Investigación CGGE, Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo del Bosque, B1900FWA, La Plata, Argentina; Consejo Nacional de Investigaciones Científicas y Técnicas (CONICET), C1033, AAJ Ciudad Autónoma de Buenos Aires, Argentina; Instituto de Astrofísica de La Plata (CCT-La Plata, CONICET-UNLP), Paseo del Bosque, B1900FWA, La Plata, Argentina", "Grupo de Investigación CGGE, Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo del Bosque, B1900FWA, La Plata, Argentina; Consejo Nacional de Investigaciones Científicas y Técnicas (CONICET), C1033, AAJ Ciudad Autónoma de Buenos Aires, Argentina; Instituto de Astrofísica de La Plata (CCT-La Plata, CONICET-UNLP), Paseo del Bosque, B1900FWA, La Plata, Argentina" ]	[ "2015MNRAS.451..791C", "2015MNRAS.452.1617H", "2016MNRAS.459.2992K", "2018A&A...616A.165V", "2019ApJ...882..132C", "2020MNRAS.498.3852B", "2023A&A...675A..90P", "2023ApJ...956..104H" ]	[ "astronomy" ]	3	[ "galaxies: dwarf", "galaxies: fundamental parameters", "galaxies: evolution", "galaxies: photometry", "galaxies: starburst", "galaxies: star formation", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1970ApJ...162L.155S", "1979MNRAS.189...95P", "1980A&A....91..269L", "1981ApJ...247..823T", "1982euse.book.....L", "1985A&AS...61...93H", "1986ApJ...300..496K", "1986ApJ...309...59L", "1990AJ....100....1F", "1994ASPC...61..327S", "1996A&AS..120..207P", "1999ApJ...510...71J", "2000A&ARv..10....1K", "2000PASJ...52..623N", "2001ApJ...553...47F", "2001ApJS..133..321C", "2002ApJ...577..164C", "2003A&A...400..451G", "2003A&A...408..929D", "2003A&A...410..481N", "2003ApJS..147...29G", "2004AJ....128.1141K", "2004PASP..116..425H", "2005A&A...429..115N", "2005A&A...429..439G", "2005AJ....130.1593V", "2005ApJS..156..345G", "2006A&A...457..477K", "2006AJ....131.1318V", "2007AJ....134..604V", "2007ApJ...663..834B", "2008A&A...477..813K", "2008A&A...487..147V", "2008MNRAS.386.1145B", "2008MNRAS.386.2311S", "2008MNRAS.391..685S", "2009A&A...501...75A", "2009ARA&A..47..481A", "2009ApJ...694..556B", "2009ApJ...707.1676C", "2010A&A...520A..90C", "2010ApJ...710..663Z", "2010ApJ...720.1738P", "2011A&A...533A..65V", "2012A&A...540A..49M", "2012MNRAS.419.2472S", "2013A&A...554A..20S", "2013A&A...559A.114M", "2013AJ....145...44Z", "2013AJ....145..101K", "2013ApJ...764...44Z", "2013MNRAS.430.1088C", "2013MNRAS.432.2731K" ]	[ "10.1051/0004-6361/201322615", "10.48550/arXiv.1403.3534" ]	1403	1403.3534_arXiv.txt	Blue compact dwarf galaxies (BCD) are part of the vast family of dwarf that also includes dwarf elliptical (dE), dwarf irregular (dI), and dwarf spheroidal (dSph) galaxies. BCDs show quite distinctive characteristics such as intense star formation at their central regions and low metallicity (e.g., Kehrig et al. \cite{keh08}; Cairos et al. \cite{cai09}), but most of them also harbor an old stellar population that accounts for most of their mass (Papaderos et al. \cite{pap96}; Cair\'os et al. \cite{cai01,cai02}; Noeske et al. \cite{noe03,noe05}; Vaduvescu et al. \cite{vad06}; Amor\'in et al. \cite{amo09}). Some BCDs carry important information related to the first stages of star formation (Sargent \& Searle \cite{sar70}; Lequeux \& Viallefond \cite{leq80}; Kunth \& Sargent \cite{kun86}) and the study of nearby BCDs located in the local universe should be used first to elucidate properties of high-redshift, low-mass galaxies that are harder to reach. According to the observational criteria proposed by Thuan \& Martin (\cite{thu81}) and Gil de Paz, Madore \& Pevunova (\cite{gil03}), BCDs have low luminosity ($M_B \gtrsim -18$ mag and $M_K \gtrsim -21$ mag) and low peak surface brightness ($\mu_{B} < 22~\rm{mag~arcsec^{-2}}$), and their integrated colours are mainly blue (mean $B-R \approx 0.7$ mag). Spectroscopically, they show intense and narrow emission lines super-imposed on a faint and blue optical continuum (Kehrig et al. \cite{keh04}). The oxygen abundance takes values, for instance, in the range $ 7.49 \leq 12 + \rm{log(O/H)} \leq 8.81$ dex in a sample of about 70 nearby field BCDs selected from the Gil de Paz et al. (\cite{gil03}) sample by Zhao, Gao \& Gu (\cite{zha10}), though abundances were obtained using different methods (see also Cairos et al. \cite{cai10}; Kehrig et al. \cite{keh13}). These results support the statement that they are metal-poor systems, since their oxygen abundances are mostly below the solar value ($12 + \rm{log(O/H)} \approx 8.7~$ dex (Asplund et al. \cite{asp09}). A correlation between the oxygen abundance and the absolute $K_S$ magnitude seems to hold for BCDs and dIs both in the Local Volume (LV), defined as the nearby Universe within $d<11$ Mpc (Karachentsev et al. \cite{kar13}), and in the nearby clusters (Vaduvescu et al. \cite{vad07,vad11}). Furthermore, the near infrared (NIR) luminosity of the old stellar population in star-forming dwarfs can be used to estimate the gas fraction (Vaduvescu et al. \cite{vad05,vad06}), defined as the gas mass over total baryonic mass (i.e., gas and stars). Vaduvescu et al. (\cite{vad11}) confirmed that the abundance - gas fraction relation for LV BCDs seems to be followed by star-forming dwarfs located in Virgo, Fornax, and Hydra clusters. Based on these results, they propose that the chemical evolution of BCDs in the Virgo and Hydra clusters seems to be consistent with the predictions of a closed-box model. They also suggest that the mass-metallicity relation followed by star-forming dwarfs looks different for environments of different densities. The dwarfs in clusters follow a mass-metallicity relation with a slope steeper than the one of the same relation for LV dwarfs. More recently, Zhao, Gao \& Gu (\cite{zha13}) have examined the oxygen abundance-gas fraction relation obtained for a sample of 53 BCDs and 22 field dIs and compared it with different models of chemical evolution. They conclude that most galaxies in their sample do not agree with a closed-box model; i.e., BCDs do not seem to have evolved as isolated systems. Since the dispersion in most of these relations is rather large, it is essential to increase the sample in order to determine the influence of the environment. During the past decades, a few possible evolutionary connections between different types of dwarfs (BCDs, dIs, dEs, dSphs) have been explored, taking into account that several physically different channels are at work, but the whole evolution scenario between them remains far from clearly understood (e.g., Kunth \& \"{O}stlin \cite{kun00}; Vaduvescu et al. \cite{vad11} and references therein; S\'anchez-Janssen et al. \cite{san13}). This paper follows Vaduvescu, McCall \& Richer (\cite{vad06}) and Vaduvescu, Richer, and McCall (\cite{vad07}) who focussed on Virgo cluster, and Vaduvescu et al. (\cite{vad11}), who focussed on the Fornax and Hydra clusters, all with the aim of studying star-forming dwarfs (BCDs and dIs) in nearby clusters and comparing their evolution with a field sample (Vaduvescu et al. \cite{vad05}; McCall et al. \cite{mcc12}). Following this investigation, in this sense we now scrutinize one small sample of star-forming dwarf candidate galaxies located in the southern cluster Antlia. The paper is organized as follows: Section~2 describes the cluster and object selection, and Section~3 presents the observations and data reduction. The results are presented in Section~4 and discussed in Section~5. Finally, Section~6 summarizes the main conclusions derived from this work.	\label{conclusions} Five star-forming dwarf galaxy candidates were selected taking from the scarce literature in the Antlia cluster with the aim of continuing our studies of physical and chemical properties of BCDs and dIs in the Local Volume and other nearby clusters (Virgo, Fornax, and Hydra). Deep $H\alpha$ and $R$ imaging was obtained for these targets in arcsec seeing Band 3 conditions using Gemini South with the GMOS-S camera. Three galaxies (FS90-98, FS90-106, FS90-147) out of five show $H\alpha$ emission and a few knots. From our GMOS spectroscopic observations, we confirm that these three galaxies are members of the Antlia cluster. We derive the oxygen abundance for these objects and find values that fall in the metallicity range derived for most of the BCDs. Shallower archival VHS $K_S$ images for all five targets were reduced and studied further, showing $K_S$ surface-brightness profile fits consistent with BCD classification, namely a sech law to account for the extended component plus a Gaussian to fit the inner outburst. FS90-155 and FS90-319 could not be confirmed as star-forming cluster candidates, and their membership status remains uncertain. FS90-155 shows some flocculent structure in both $R$ and $K_S$ images resembling spiral arm structures and showing no $H\alpha$ emission. FS90-319 presents a flat flux distribution in its inner region that has one non-central diffuse emission visible in both visible and NIR and no $H\alpha$ emission, so their membership status remains uncertain. Two-dimensional physical and chemical relations were studied using $K_S$ sech data for the whole Antlia sample to probe wheather known relations (specifically the size-luminosity and luminosity-metallicity) are followed by dwarfs in the LV and the Virgo, Fornax and Hydra clusters. Future studies could benefit from deeper NIR imaging and radio data in the aim to compare star-forming dwarf formation and galaxy evolution in the nearby Universe.	14	3	1403.3534	Context. The formation and evolution of dwarf galaxies in clusters need to be understood, and this requires large aperture telescopes. <BR /> Aims: In this sense, we selected the Antlia cluster to continue our previous work in the Virgo, Fornax, and Hydra clusters and in the Local Volume (LV). Because of the scarce available literature data, we selected a small sample of five blue compact dwarf (BCD) candidates in Antlia for observation. <BR /> Methods: Using the Gemini South and GMOS camera, we acquired the Hα imaging needed to detect star-forming regions in this sample. With the long-slit spectroscopic data of the brightest seven knots detected in three BCD candidates, we derived their basic chemical properties. Using archival VISTA VHS survey images, we derived K<SUB>S</SUB> magnitudes and surface brightness profile fits for the whole sample to assess basic physical properties. <BR /> Results: FS90-98, FS90-106, and FS90-147 are confirmed as BCDs and cluster members, based on their morphology, K<SUB>S</SUB> surface photometry, oxygen abundance, and velocity redshift. FS90-155 and FS90-319 did not show any Hα emission, and they could not be confirmed as dwarf cluster star-forming galaxies. Based on our data, we studied some fundamental relations to compare star forming dwarfs (BCDs and dIs) in the LV and in the Virgo, Fornax, Hydra, and Antlia clusters. <BR /> Conclusions: Star-forming dwarfs in nearby clusters appear to follow same fundamental relations in the near infrared with similar objects in the LV, specifically the size-luminosity and the metallicity-luminosity, while other more fundamental relations could not be checked in Antlia due to lack of data. <P />Based on observations acquired at Gemini South (GS-2010A-Q-51 and GS-2012A-Q-59) and ESO VISTA Hemisphere Survey (VHS).	false	[ "Antlia clusters", "Hydra clusters", "basic physical properties", "star forming dwarfs", "nearby clusters", "same fundamental relations", "cluster members", "clusters", "dwarf galaxies", "LV", "dwarf cluster star-forming galaxies", "archival VISTA VHS survey images", "large aperture telescopes", "Antlia", "ESO VISTA Hemisphere Survey", "data", "surface brightness profile", "Hydra", "Fornax", "VHS" ]	10.689938	7.532768	-1
483244	[ "Kuhn, Michael A.", "Feigelson, Eric D.", "Getman, Konstantin V.", "Baddeley, Adrian J.", "Broos, Patrick S.", "Sills, Alison", "Bate, Matthew R.", "Povich, Matthew S.", "Luhman, Kevin L.", "Busk, Heather A.", "Naylor, Tim", "King, Robert R." ]	2014ApJ...787..107K	[ "The Spatial Structure of Young Stellar Clusters. I. Subclusters" ]	130	[ "Department of Astronomy & Astrophysics, Pennsylvania State University, 525 Davey Laboratory, University Park, PA 16802, USA", "Department of Astronomy & Astrophysics, Pennsylvania State University, 525 Davey Laboratory, University Park, PA 16802, USA", "Department of Astronomy & Astrophysics, Pennsylvania State University, 525 Davey Laboratory, University Park, PA 16802, USA", "School of Mathematics and Statistics, University of Western Australia, 35 Stirling Highway, Crawley, WA 6009, Australia", "Department of Astronomy & Astrophysics, Pennsylvania State University, 525 Davey Laboratory, University Park, PA 16802, USA", "Department of Physics, McMaster University, 1280 Main Street West, Hamilton, ON L8S 4M1, Canada", "Department of Physics and Astronomy, University of Exeter, Stocker Road, Exeter, Devon EX4 4SB, UK", "Department of Physics and Astronomy, California State Polytechnic University, 3801 West Temple Avenue, Pomona, CA 91768, USA", "Department of Astronomy & Astrophysics, Pennsylvania State University, 525 Davey Laboratory, University Park, PA 16802, USA", "Department of Astronomy & Astrophysics, Pennsylvania State University, 525 Davey Laboratory, University Park, PA 16802, USA", "Department of Physics and Astronomy, University of Exeter, Stocker Road, Exeter, Devon EX4 4SB, UK", "Department of Physics and Astronomy, University of Exeter, Stocker Road, Exeter, Devon EX4 4SB, UK" ]	[ "2014ApJ...787..108G", "2014ApJ...787..109G", "2014ApJ...794..147P", "2014ApJ...795...55D", "2015A&A...573A..19S", "2015A&A...573A..95M", "2015A&A...578A..35M", "2015AJ....150...78Z", "2015ApJ...798..126J", "2015ApJ...802...60K", "2015ApJ...811...10R", "2015ApJ...812..131K", "2015MNRAS.446..226L", "2015MNRAS.447..728B", "2015MNRAS.453.1026K", "2015arXiv150704313H", "2016A&A...586A..68P", "2016AJ....151....5M", "2016AJ....151..126S", "2016ApJ...819..139B", "2016ApJ...821...98K", "2016ApJ...827...52L", "2016ApJ...827...96F", "2016ApJ...833..193R", "2016IAUTA..29..502C", "2016MNRAS.458.3027M", "2016MNRAS.461.2519V", "2016MNRAS.461.2953H", "2017A&A...607A..86R", "2017AJ....153..122Z", "2017AJ....154...87K", "2017AJ....154..214K", "2017ApJ...836...98J", "2017ApJ...838...61P", "2017ApJS..229...28G", "2017IAUS..316...55Z", "2017MNRAS.471.3699M", "2017MNRAS.472.2808D", "2017MNRAS.472.4982S", "2017arXiv171111101K", "2018A&A...609A..10V", "2018A&A...620A..27J", "2018AJ....156...84K", "2018ASSL..424....1A", "2018ASSL..424..119F", "2018ASSL..424..143B", "2018ApJ...858...31S", "2018ApJ...860L...4S", "2018ComAC...5....2V", "2018MNRAS.476.1213G", "2018MNRAS.477..298G", "2018MNRAS.477.1903S", "2018MNRAS.477.5191R", "2018MNRAS.481.1679P", "2018PASP..130g2001G", "2018PhRvE..98d2133M", "2019A&A...622A.118S", "2019A&A...622A.184B", "2019A&A...623A..25D", "2019A&A...623A.159P", "2019A&A...624A..34Z", "2019A&A...625A.134R", "2019AJ....158..235G", "2019ARA&A..57..227K", "2019ApJ...870...32K", "2019ApJ...871...38K", "2019ApJ...871...46K", "2019ApJ...878..111H", "2019ApJ...881...37P", "2019ApJ...881...79M", "2019ApJ...884....6M", "2019ApJ...884..173K", "2019JPhCS1231a2028D", "2019MNRAS.486.2477W", "2019MNRAS.486.3019F", "2019MNRAS.486.4354R", "2019MNRAS.487.2977G", "2020A&A...636A..80B", "2020A&A...640A..84D", "2020A&A...640A..85D", "2020A&A...642A..21R", "2020A&A...644A.141R", "2020ApJ...891...81P", "2020IAUS..351..357K", "2020MNRAS.493.4925D", "2020MNRAS.499..748D", "2020SSRv..216...69A", "2021A&A...645A..84M", "2021A&A...645A..94N", "2021A&A...647A..14G", "2021A&A...656A..49H", "2021ApJ...913...95P", "2021ApJ...923..129J", "2021MNRAS.506.3239G", "2021MNRAS.506.4603D", "2021MNRAS.506.5781L", "2022A&A...659A.169D", "2022A&A...660A..61D", "2022A&A...668A..19M", "2022AJ....163..266L", "2022ARep...66..361V", "2022ApJ...926..141T", "2022ApJ...935..142H", "2022ApJ...937...46K", "2022ApJ...939...34H", "2022MNRAS.512.2584L", "2022MNRAS.514..920D", "2022MNRAS.515..167G", "2022MNRAS.516.5258S", "2022MNRAS.517..161K", "2022PASP..134d2001M", "2022arXiv220310007W", "2022csss.confE..50M", "2023A&A...669A..22P", "2023AJ....165....3K", "2023AJ....166...97L", "2023AJ....166..183V", "2023ASPC..534....1C", "2023ASPC..534..129W", "2023ApJ...944..211L", "2023MNRAS.519.3643B", "2023MNRAS.521.1338C", "2023arXiv230802279C", "2023arXiv231108358W", "2024AJ....167..230C", "2024AstBu..79...71D", "2024BSRSL..93..582V", "2024MNRAS.529.3925B", "2024MNRAS.530.4970Y", "2024arXiv240610388V" ]	[ "astronomy" ]	28	[ "H II regions", "ISM: structure", "methods: statistical", "open clusters and associations: general", "stars: formation", "stars: pre-main sequence", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1922RSPTA.222..309F", "1942psd..book.....C", "1962AJ.....67..471K", "1967ApJ...147..799B", "1967ApJ...149L...1K", "1973ApJ...185L.131K", "1996MNRAS.279.1037G", "1998AJ....116..322H", "1998ApJ...492..540H", "2000ApJ...530..277E", "2000prpl.conf..151C", "2001AJ....121.1507E", "2003ARA&A..41...57L", "2003MNRAS.338..120M", "2003MNRAS.339..577B", "2003PASP..115..953B", "2004AJ....127.1014S", "2004AJ....128.1254L", "2004ApJ...616..288B", "2004JGRB..109.3308O", "2004MNRAS.348..589C", "2005AJ....130..134B", "2005ApJS..160..319G", "2005ApJS..160..379F", "2005ApJS..160..557S", "2005ApJS..160..582P", "2006AJ....131.1163S", "2006ApJ...641..389R", "2006ApJ...641L.121T", "2006MNRAS.373..752G", "2007A&A...468..353G", "2007A&A...476..199A", "2007ASPC..371...22B", "2007ApJ...655L..45M", "2007sptz.prop40791M", "2008ApJ...672..861R", "2008ApJ...675..464W", "2008ApJ...676.1109F", "2008MNRAS.383..375C", "2008MNRAS.386.1855C", "2008gady.book.....B", "2009A&A...498L..37P", "2009A&ARv..17..309G", "2009AJ....138..227F", "2009ApJ...699.1454G", "2009ApJ...700L..99A", "2009ApJ...704..453F", "2009ApJS..181..321E", "2009ApJS..184...18G", "2009MNRAS.392..590B", "2009MNRAS.395.1449A", "2009yCat....1.2023S", "2010A&A...518L.102A", "2010ARA&A..48..431P", "2010ApJ...708.1760G", "2010ApJ...723L...7K", "2010ApJ...725.2485K", "2010MNRAS.404..721M", "2010MNRAS.404.1061M", "2010MNRAS.409..628H", "2010MNRAS.409L..54B", "2010PNAS..107.7153F", "2011A&A...535A.128B", "2011A&A...536A..90P", "2011AN....332..172S", "2011ApJ...727...88C", "2011ApJS..194....2B", "2011ApJS..194....3G", "2011ApJS..194....9F", "2011ApJS..194...11W", "2011ApJS..194...14P", "2011ApJS..194...16T", "2011MNRAS.410.2339B", "2011MNRAS.416..383S", "2012A&A...540L..11S", "2012A&A...543A...8M", "2012A&A...545A.122P", "2012A&A...547A..97A", "2012AJ....144..192M", "2012MNRAS.426.2917G", "2012MNRAS.426L..11G", "2012MNRAS.427..637P", "2013A&A...555A.135P", "2013ApJ...764...29B", "2013ApJ...769..140Y", "2013ApJS..209...26F", "2013ApJS..209...27K", "2013ApJS..209...28K", "2013ApJS..209...29K", "2013ApJS..209...30N", "2013ApJS..209...31P", "2013ApJS..209...32B", "2014ApJ...787..108G", "2014ApJ...787..109G", "2014ApJS..213....1T" ]	[ "10.1088/0004-637X/787/2/107", "10.48550/arXiv.1403.4252" ]	1403	1403.4252_arXiv.txt	} Massive star-forming regions (MSFRs) are a major, and perhaps the dominant, mode of star formation in the Galaxy \citep{Lada03,Fall09,Chandar11}. Most of the young stars in these massive complexes are clustered \citep{Clarke00,Lada03}, but a substantial, spatially distributed population of young stars also exists \citep{Evans09,Feigelson11}. Open questions on how young stellar clusters form and evolve in time in these harsh environments are outlined in the overview of the Massive Young Star-Forming Complex Study in Infrared and X-ray \citep[\mystix][]{overview}. This multifaceted project starts by characterizing the young stellar populations of 20 OB-dominated MSFRs at distances $d<3.6$~kpc. The \mystix\ analysis provides a large sample of 31,754 young stellar members of these MSFRs \citep{mpcm}. The spatial distributions of young stars will show the imprint of cluster formation and evolution. For example, \citet{Elmegreen00} suggests that star formation occurs in a crossing time in a cloud with freely decaying turbulence, leaving behind stars that trace the subclustered structure of the natal cloud. However, \citet{Tan06} argue that gradual star formation in clouds with continually driven turbulence is necessary to explain smooth stellar surface-density gradients in clusters like the Orion Nebula Cluster (ONC). In addition, numerical simulations hint that dynamical processes of merging stellar subclusters may have a significant effect on the properties of the young clusters that are produced \citep[e.g.][]{McMillan07,Bate09a,Maschberger10}. Observations of young stellar cluster sizes, densities, and morphologies will enable testing the theoretical models of star formation. However, the interpretation of spatial distributions faces several challenges. \begin{enumerate} \item Young stellar clusters, both real and simulated, form with a wide variety of intricate morphologies. Therefore, summary statistics must be developed which capture the salient and astrophysically meaningful features of the distribution of star positions so that structure in one region may be compared quantitatively to another \citep[e.g.][]{Cartwright04,Allison09b,Gutermuth09}. For example, recent cluster-formation simulations produce stellar spatial distributions that resemble (to the eye) the stars in real star-forming regions \citep[e.g.][]{Bate09a,Bonnell11}, but robust comparison of structure should be based on calibrated statistical measures. \item There may be ambiguities in determining which stars are clustered \citep{Bressert10,Gieles12,Pfalzner12}. Furthermore, molecular clouds and star formation show signs of fractal structure ranging from Galactic scales \citep{Elmegreen01} down to subcluster scales \citep{Schmeja11}. Were stellar distributions truly scale-invariant, the determination of cluster size and density would depend on the arbitrary choice of cluster boundaries. \item Different mechanisms of star cluster formation may produce clusters with similar properties. For example, a young stellar cluster that has relaxed through two-body interactions may appear similar to a cluster that underwent rapid, violent relaxation induced by the merger of subclusters \citep{Allison09a} or gas removal \citep{Moeckel10}. \end{enumerate} A particularly salient features of the spatial distributions of young stars is the organization in clusters and subclusters\footnote{ The definitions of ``young stellar cluster'' in the star-formation literature are divergent \citep[for an alternative definition cf.][]{PortegiesZwart10}. Here, we use ``cluster'' in a statistical sense (subclusters are the components of the finite mixture model in Section 3). }. \citet{Hillenbrand98} found that a \citet{King62} profile describes the radial surface density profile of the ONC, with $\rho_0=2-3\times 10^4$~M$_\odot$~pc$^{-3}$ and core radius $r_0=0.15-0.2$~pc. This profile has also been fit to NGC~6611, W~40, Tr~15, NGC~3603, the Arches cluster, h and $\chi$ Per, and young stellar clusters in the Large Magellanic Clouds \citep{Wang08, Kuhn10, Wang11, Sung04, Harfst10, Bragg05,Mackey03}. \citet{Pfalzner12} note that these clusters tend to have large ratios of half-mass radius to core radius, so the truncated King profile is a better model than the Plummer sphere. In addition to the main subclusters in regions, smaller groups of stars can be seen outside the main clusters, or even as subclusters within the main clusters. It is unclear whether a relaxed surface density profile is a good model for subclusters, in part because young stellar clusters are generally not expected to have had time to dynamically relax through two-body interactions and molecular gas may contribute significantly to the gravitational potential. Nevertheless, similar surface density profiles have been used by \citet{Smith11}, who use Plummer sphere subclusters in their simulations, and found by \citet{Maschberger10} in the radial density profile of simulated subcluster mergers. We use the young stellar samples from the \mystix\ project to analyze the subcluster properties in 17 MSFRs. This rich statistical sample of stars is particularly useful for investigation of spatial structure, since, prior to \mystix, the stellar populations in most of these regions were poorly characterized. \mystix\ Probable Complex Members include low-mass and high-mass stars and disk-bearing and disk-free stars \citep{mpcm}, and the multiwavelength X-ray/infrared observations ameliorate effects of high absorption from the natal cloud, infrared (IR) and optical nebulosity from ionized gas, and source crowding \citep{overview}. In particular \mystix\ greatly improves the detection of stars at the centers of dense subclusters, where IR-only studies often fare the worst \citep[e.g.][]{Bressert10}. Although the \mystix\ samples are not ``complete,'' the identification of young stars is performed in a uniform way for the different regions which allows for comparative analysis of stellar populations in these different regions. Here, we analyze the projected spatial distributions of young stars from these censuses; analysis is done in parsec units to compare regions at different distances, and the distribution of star positions is analyzed using statistical theory for spatial point processes. This work (Paper I) discusses subclusters in these regions. The subclusters have a variety of configurations that relate to different ways that star formation can occur and different stages in the evolution of a MSFR. Paper II will analyze the intrinsic stellar populations of individual subclusters cataloged here by correcting for incompleteness in each region. We recognize that the full spatial substructure of stars in a star-forming region cannot be reduced to discrete subclusters if fractal clustering is present; this will be investigated in Paper III using non-parametric, parametric, and stochastic methods for characterizing spatial structures. Mass segregation, the concentration of massive stars in the cores of many clusters will also be investigated. Other papers in the \mystix\ project will use the subclusters derived here including study of the stellar ages \citep{Getman13a}, protoplanetary disk fractions, and spatial relationships between subclusters and their cloud environs. This paper is organized as follows. In Section \ref{sample_section}, the \mystix\ census of young stars is described, and biases associated with the multiwavelength data are discussed. In Section~\ref{methodology_section} the statistical method for identifying and modeling subclusters is described. The cluster fitting results and catalogs are presented in Section \ref{results_section}. The structure of star-forming complexes are examined in Section~5, and subcluster properties are investigated in Section~6. Section~7 is the discussion. An Appendix gives code permitting the reader to recover and revise the subclusters obtained here using published \mystix\ source tables.	} The MPCM samples have been designed to improve the statistical characterization of stellar populations in MSFRs \citep{overview}. Using a rich statistical sample of 16,608 MPCMs, ``flattened'' to remove X-ray spatial sensitivity variations, and the finite-mixture-model technique to identify subclusters among these young stars, new details about the spatial structure of young stellar clusters emerge. In the diverse MSFRs surveyed by the \mystix\ project, we see a picture of star-forming regions with multiple non-coeval subclusters that probably formed with similar small sizes, but have expanded to a wide range of sizes. A similar finding of cluster expansion in a different sample is studied by \citet{Pfalzner09}. The catalog of 142 \mystix\ subclusters derived here, combined with inferred intrinsic stellar populations (Paper~II) and age estimates for the subclusters \citep{Getman13a}, will help provide powerful links between observational and theoretical analysis star-cluster formation. The finite mixture model, using isothermal ellipsoids, is a novel cluster analysis method for MSFRs that can reveal information that is hard to obtain from other cluster-finding methods. \citet{Pfalzner11} finds that multimodal properties in regions with varying cluster surface density and richness are difficult to determine from surface density distributions alone. But our method is an objective and mathematically founded procedure to separate multiply-clustered patterns into distinct isothermal ellipsoidal subclusters. The procedure is flexible, finding statistically significant structures on all scales simultaneously, even treating cases such as small clusters reside inside large ones. This is important because, even in relatively simple MSFRs like the Orion Nebula, layered subcluster structures are present and can affect inferred subcluster properties. Furthermore, our algorithm does not require a single threshold in surface-density or separation to divide subclusters, which can help avoid bias when subclusters with very different properties exist in the same star-forming region. However, the finite-mixture-model method requires additional assumptions about region structure and subcluster shape. We validate the subcluster results using residual maps, $\chi^2$ tests of source counts, and comparison with results from previous studies of these regions. The models reproduce the observed surface densities reasonably well, although residual maps reveal some extra structure. The improved stellar census in dense regions improves the accuracy of subcluster radius and peak surface density maps. On larger scales, we caution that the MPCM sample coverage of widely distributed young stars is limited, and isothermal ellipsoid fitting is not sensitive to any truncation of cluster sizes that may be present. \subsection{The Structure of Stellar Populations in MSFRs} The numbers of subclusters and their arrangements vary significantly from region to region. Heuristically, the clustering morphologies are divided into four classes: isolated clusters, core-halo clusters, clumpy clusters, and linear chains of clusters (Section~\ref{classes_section}). The dominant subcluster in a region (i.e.\ the subcluster with most stars, $N_{4,\mathrm{obs}}$) often sits at the center of a group of subclusters (e.g.\ Subcluster L in M~17), but this is not always true (e.g.\ subcluster B in the Eagle Nebula). In regions with linear chain structure there may be no subcluster that is particularly dominant. While, in contrast, regions with simple or core-halo structure may be much more centrally concentrated. In DR~21, NGC~2264, and NGC~6334, the linear chain cluster structure maps onto molecular filaments in the regions, and it is clear that the star cluster structure was at least partially inherited from the filamentary chain of molecular cloud clumps. More detail about correspondence between subcluster locations and molecular cloud clump locations will be given in a later \mystix\ study. Thus, regions with linear chain structure have not evolved much from their initial state at the time of the birth of stars. The more centrally concentrated clumpy, core-halo, and simple cluster structures are more likely to have dynamically evolved and, in cases where the subclusters are well fit by single isothermal ellipsoids, may have achieved some dynamical relaxation. The morphological classification suggests an evolutionary progression from linear chain structure to clumpy structure to core-halo structure to simple cluster structure. Hierarchical cluster formation, often in molecular filaments, would evolve into dynamical relaxed unimodal structures. However, the \mystix\ regions indicate that this progression is not always followed. For example, the linear chain region NGC~1893 is almost entirely unembedded, and the chain of subclusters in NGC~6334 exhibit a wide range of absorptions. The distribution of subcluster properties show trends common to all the regions. A lognormal distribution of subcluster size peaks at 0.17~pc, and many subclusters have ellipticities of $\epsilon \leq 0.3$, but few have ellipticities greater than 0.5. No relationships between ellipticity and size are seen. An additional property of subclusters, their absorption, can help link these properties to cluster formation history and evolution. Line-of-sight absorption is measured both by the near-IR color index $J-H$ and the X-ray Median Energy indicator \citep{Getman10}. While absorption may be caused by overlaying molecular material, in most cases it represents local cloud cores within which the youngest clusters reside. In complex \mystix\ star-forming regions, a wide range of absorption can be seen for the different subclusters, indicating these subclusters are not coeval. A statistically significant anti-correlation between absorption and size of a subcluster is found, most easily explained if both absorption and size are related to age. This issue is pursued in the accompanying study by \citep{Getman13a} where spatial-age gradients in \mystix\ regions are found based on a new age indicator for individual pre-main sequence stars. A similar relationship was found for subclusters within the Rosette Molecular Cloud by \citep{Ybarra13} who infer a gas removal $e$-folding timescale of 0.4~Myr. Cluster expansion is expected to produce a wide range of stellar surface densities seen in many star-forming regions \citep{Gieles12}. Early cluster expansion arises principally from the loss of the gravitational potential of the molecular material, but stellar dynamics can also produce stellar halos around dense cores through three-body interactions with hard binaries and gravothermal core collapse associated with mass segregation \citep{Giersz96}. This issue will be discussed further in Paper II when intrinsic populations, rather than observed population subject to different sensitivity effects, are derived. \subsection{Links to Astrophysical Theory of Cluster Formation} The anti-correlation between radius and absorption provides evidence that subcluster expansion is important part of subcluster dynamical evolution. From numerical simulations, \citet{Moeckel10} find that even the very dense simulated cluster produced by \citet{Bate09a}, with a half-mass radius of $\sim$0.05~pc, grew to $>$2~pc over a period of $\sim$2~Myr. The typical small size that we find for the embedded clusters indicates that the high initial densities from the simulation in \citet{Bate09a} are realistic, which can make competitive accretion and dynamical interaction between protostars much more important than if clusters form $\sim$1~pc in size. Simulation of cluster expansion after gas expulsion by \citet{Goodwin06} show a relationship between star-formation efficiency and the age-radius relationship. They find that, after 2.5~Myr, core radii will grow by a factor of $\sim$1.5 for a star-forming efficiency of 60\% and a factor of $\sim$3 for a star-forming efficiencies of 10--30\%. Larger expansion factors are found in models of very rich clusters by \citet{Banerjee13}. For comparison, \mystix\ results indicate that subcluster core radii start out at $\sim$0.08~pc and grow to the typical $\sim$0.2~pc for the typical unabsorbed \mystix\ subclusters (a factor of more than 2 expansion). The stars in \mystix\ regions are generally $<$5~Myr old, but more precise age determination is necessary to determine which age-radius relation provides the better fit. \citet{Pfalzner11} find a trend consistent with a star-formation efficiency of $\sim$30\% using the clusters in \citet{Lada03}. The ellipticities of \mystix\ subclusters could be inherited from the structure of the parental molecular cloud or could arise from subcluster mergers. The subcluster ellipticity distribution resembles those found by numerical simulations of merging subclusters \citep[][their Figure~10b]{Maschberger10} in which merger products have a distribution of ellipticities, including a tail at high ellipticities. The cluster morphological classes may correspond to structures seen in star formation simulations, such as those performed by \citet{Bate03,Bate09a} or \citet{Parker12}. Simulations often show multiple pockets of star formation occurring along gas filaments as the cloud collapses, which can end up merging and dynamically relaxing as the simulation progresses. Clumpy clusters may indicate early stages of partial merging, core-halo structures may be an intermediate stages of mergers, and isolated clusters with relaxed structure a final stage. \citet{Getman13b} add an important constraint to such a scenario: in at least two \mystix\ regions, the dense core of the cluster has younger stars than the outer regions. This observational result requires some special behavior, such as continual feeding of gas towards the cluster center or dispersal of older core stars into the halo. The simple clusters may represent those that have completed subcluster merging and dynamical relaxation. Subcluster morphology can also influence these rates. \citet{Smith11} find that the cluster formation rate is related to a filling factor measuring the fraction of volume containing subclusters. When this filling factor becomes high, tidal interactions accelerate subcluster mergers and formation of a simple unimodal cluster is promoted. This phenomenon may be occurring today in crowded, clumpy \mystix\ regions like M~17. The core-halo structures in 5 out of 17 \mystix\ regions may also be a result of subcluster mergers. \citet{Bate09a} report a hydrodynamical and N-body simulation of a MSFR resulting in a cluster with a half-mass radius of 0.05~pc and a halo of stars extending out $\sim$0.4~pc. \citet{Smith11} found that gravitational interactions between subclusters of young stars in simulated star-forming environments can scatter stars out of the clumps, which can lead to a dense core surrounded by a less dense halo of scattered stars. \citet{Maschberger10} also identify a halo of scattered low-mass stars around the dominant cluster, albeit with a surface-density--radius relation opposite of the effect we find with concentric ellipsoids. However, \citet{Pfalzner13} report that 2-body encounters can produce a population of cluster stars with high ellipticities which will spend most of the time in a halo. The accompanying paper by \citet{Getman13b} describe star formation scenarios that may account for core-halo structures. All of these issues relating to cluster formation and MSFR star formation histories will be pursued in \mystix\ papers by \citet{Getman13a} and Paper II where quantitative measures of subcluster intrinsic populations and cluster ages will be combined with the structural results obtained here.	14	3	1403.4252	The clusters of young stars in massive star-forming regions show a wide range of sizes, morphologies, and numbers of stars. Their highly subclustered structures are revealed by the MYStIX project's sample of 31,754 young stars in nearby sites of star formation (regions at distances <3.6 kpc that contain at least one O-type star.) In 17 of the regions surveyed by MYStIX, we identify subclusters of young stars using finite mixture models—collections of isothermal ellipsoids that model individual subclusters. Maximum likelihood estimation is used to estimate the model parameters, and the Akaike Information Criterion is used to determine the number of subclusters. This procedure often successfully finds famous subclusters, such as the BN/KL complex behind the Orion Nebula Cluster and the KW-object complex in M 17. A catalog of 142 subclusters is presented, with 1-20 subclusters per region. The subcluster core radius distribution for this sample is peaked at 0.17 pc with a standard deviation of 0.43 dex, and subcluster core radius is negatively correlated with gas/dust absorption of the stars—a possible age effect. Based on the morphological arrangements of subclusters, we identify four classes of spatial structure: long chains of subclusters, clumpy structures, isolated clusters with a core-halo structure, and isolated clusters well fit by a single isothermal ellipsoid.	false	[ "young stars", "star formation", "stars", "subclusters", "famous subclusters", "massive star-forming regions", "finite mixture models", "subcluster core radius", "region", "regions", "clumpy structures", "spatial structure", "isolated clusters", "The subcluster core radius distribution", "31,754 young stars", "numbers", "that model individual subclusters", "lt;3.6 kpc", "isothermal ellipsoids", "nearby sites" ]	10.061704	9.287377	-1
482915	[ "Anderson, Michael E.", "Bregman, Joel N." ]	2014ApJ...785...67A	[ "Modeling X-Ray Emission around Galaxies" ]	9	[ "Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA; Max-Planck Institut für Astrophysik, Karl-Schwarzschild-Strasse 1, D-85748 Garching bei München, Germany;", "Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA" ]	[ "2014MNRAS.445..175G", "2016MNRAS.455..227A", "2016MNRAS.458.3773S", "2016SPIE.9905E..4PB", "2017ApJ...834..164B", "2018ApJ...862....3B", "2018ApJ...867...14W", "2019ApJ...877...91B", "2020ApJ...903...35H" ]	[ "astronomy" ]	4	[ "galaxies: halos", "galaxies: individual: NGC 720", "methods: statistical", "X-rays: diffuse background", "X-rays: galaxies", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Astrophysics of Galaxies" ]	[ "1976A&A....49..137C", "1977ARA&A..15..541G", "1979ApJ...234L..27F", "1982ARA&A..20..547F", "1985ApJ...293..102F", "1989ARA&A..27...87F", "1990ApJ...354..211S", "1993ApJ...404L...9M", "1994ApJ...436...44E", "1995A&A...296....1C", "1996MNRAS.283..690P", "1998ApJ...493..715S", "2000ARA&A..38..289M", "2001ApJ...562..575D", "2001MNRAS.328..461O", "2002A&A...389...93L", "2002ARA&A..40..539R", "2003ARA&A..41..191M", "2003ASPC..295..477C", "2003ApJ...583...70M", "2003ApJ...598..886U", "2003MNRAS.344L..13E", "2004AJ....128.2048B", "2004ApJS..151..193S", "2004MNRAS.355..694M", "2005A&A...436...67S", "2006A&A...448...43T", "2006ApJ...646..899H", "2006MNRAS.371..147L", "2006NewA...11..465P", "2007ApJ...661L.117H", "2007ApJ...666..147G", "2007ApJ...671.1523H", "2008A&A...478..575K", "2008ApJ...684.1075M", "2009A&A...496..343E", "2009A&A...501..899R", "2009ApJ...695.1127G", "2009ApJ...697...79R", "2009ApJ...703..982G", "2009MNRAS.394.1741O", "2009MNRAS.395..657G", "2009PASJ...61.1117B", "2009PASJ...61S.291Y", "2010ApJ...714..423K", "2010ApJ...715L...1M", "2010ApJ...719..119D", "2010ApJS..189...37E", "2010CAMCS...5...65G", "2010MNRAS.407..263A", "2010PASJ...62..371H", "2011A&A...529A.133E", "2011ApJ...729...53H", "2011ApJ...737...22A", "2011ApJS..195...10X", "2011PASJ...63S1019A", "2011Sci...331.1576S", "2012A&A...541A..57E", "2012ApJ...752...46L", "2012ApJ...755..107D", "2012ApJ...758L..41T", "2012MNRAS.422.3503W", "2012MNRAS.424.1826W", "2012MNRAS.425.2069M", "2012NJPh...14b5010B", "2012PASJ...64...95S", "2013ApJ...762...20F", "2013ApJ...762..106A", "2013ApJ...765..141N", "2013ApJ...766...90I", "2013ApJ...772...97B", "2013ApJS..208...19H", "2013MNRAS.428.2085L", "2013MNRAS.429.3288S", "2013MNRAS.432..554W", "2013PASP..125..306F", "2013arXiv1302.5086R", "2014MNRAS.437.3939U" ]	[ "10.1088/0004-637X/785/1/67", "10.48550/arXiv.1403.0584" ]	1403	1403.0584_arXiv.txt	The study of very extended emission comprises a large portion of the work of extragalactic X-ray astronomy. All galaxy clusters (\citealt{Gursky1977}, \citealt{Forman1982}, \citealt{Rosati2002}) and most galaxy groups (\citealt{Mulchaey1993}, \citealt{Ebeling1994}, \citealt{Ponman1996}, \citealt{Mulchaey2000}) are suffused with a hot ($kT > 10^6$ K), X-ray emitting gaseous medium. In all clusters and many groups, this medium contains the majority of the baryons in the system (\citealt{Ettori2003}, \citealt{Gonzalez2007}, \citealt{Giodini2009}, \citealt{Andreon2010}, \citealt{Dai2010}, \citealt{Sanderson2013}), and extends to hundreds of kpc. In recent years, X-ray observations have even been able to push outwards to the virial radius of some nearby clusters (e.g. \citealt{George2009}, \citealt{Bautz2009}, \citealt{Reiprich2009}, \citealt{Hoshino2010}, \citealt{Kawaharada2010}, \citealt{Simionescu2011}, \citealt{Akamatsu2011}, \citealt{Bonamente2012}, \citealt{Morandi2012}, \citealt{Sato2012}, \citealt{Walker2012a}, \citealt{Walker2012b}, \citealt{Ichikawa2013}, \citealt{Walker2013}, \citealt{Urban2014}). Individual galaxies are also surrounded by extended X-ray emitting halos. Around elliptical galaxies, these hot gaseous halos have been studied for decades (e.g. \citealt{Forman1979}, \citealt{Forman1985}, \citealt{Fabbiano1989}, \citealt{Mathews2003}, \citealt{O'sullivan2001}, \citealt{Mulchaey2010}). Starbursting spirals have extended coronae above and below the disk extending to a few tens of kpc (\citealt{Strickland2004a}, \citealt{Li2006}, \citealt{Tullmann2006}, \citealt{Owen2009}, \citealt{Yamasaki2009}, \citealt{Li2013}). We also recently reported the detection of hot gaseous halos around more quiescent massive spiral galaxies, extending out to $\sim 50$ kpc (\citealt{Anderson2011}, \citealt{Dai2012}, \citealt{Anderson2013}); \citet{Bogdan2013} have confirmed one of these detections and have discovered another hot halo as well. Very extended hot gas is also detected around merging galaxies such as NGC 6240 \citep{Nardini2013}. As both of these fields continue to detect emission at larger radii and lower X-ray surface brightness, it is becoming increasingly important to have effective observational techniques for studying faint, extended X-ray emission. At present, there are two major approaches to this analysis: spectral fitting and spatial binning. Spatial binning is conceptually simple: one measures the X-ray radial surface brightness profile in a given band, and infers a gas density profile from the surface brightness profile. The major uncertainties with this method are flat-fielding the image and estimating the background. For bright sources, blank-sky backgrounds are sufficient, but for faint emission the background should be estimated in-field, which requires accurate flat-fielding. A secondary concern is separating various components of emission, if they exist; examples of this separation are illustrated in \citet{Anderson2011}, \citet{Dai2012}, and \citet{Bogdan2013}. For the spectral method, one measures the X-ray spectrum in various regions, instead of the surface brightness profile. Analyzing a spectrum requires more photons than measuring a broad-band surface brightness, in part because most realistic spectra have many more free parameters than a surface brightness profile. Thus, the image is usually broken into large radial annuli, which sacrifices some location information. On the other hand, the various instrumental and background components are included in the spectral model, so in theory this method does not require separate flat-fielding or background subtraction. Also additional source components can be included, obviating concerns about confusion between hot gas emission and other X-ray sources such as X-ray binaries or background point sources. The primary downsides to this method are model specification and the need for more photons. The aforementioned issues with these methods can be quite important, and can lead to conflicting results. For the spatial method, one major failure mode in flat-fielding is incorrect estimation of the ``vignetting'' - the decrease in sensitivity of the detectors as off-axis angle increases. The vignetting profile varies as a function of time (typically getting worse as the telescope degrades), and must be computed separately for each observation. It also varies as a function of energy, so a given exposure map (which contains the information about the vignetting) should only be computed for and applied to a narrow energy band over which the vignetting effects do not vary significantly. An example of this difficulty occurs in \citet{Pedersen2006}, where the authors found evidence for a hot halo around the giant spiral NGC 5746 at $4\sigma$ significance. Later, after applying an updated calibration file which accounted more correctly for the time-dependent degradation of the instruments, the vignetting profile changed and the signal disappeared \citep{Rasmussen2009}. For the spectral method, model specification is particularly important. In the observations of interest today, the signal from hot gas is lower than the background, so the model for the background components in the spectrum can significantly influence the inferred properties of the signal. It is not trivial to construct a model for the X-ray background, and most studies use slightly different prescriptions. The components generally include the extragalactic AGN background at hard energies, and Solar wind X-rays (time variable), emission from the Local Hot Bubble, and the Galactic hot halo at softer energies. These components are all variable (spatially and sometimes temporally) and so their normalizations are not typically known a priori. Moreover, the emission itself can be quite spectrally complex - especially for the Local Bubble and Solar wind X-rays, where charge exchange can affect the signal and the gas is not necessarily in collisional equilibrium. Without sufficient photons to fit all of these spectral components, this method is very susceptible to systematic errors, either from degeneracies between free model parameters, or as a result of fitting the data with an inappropriate model and/or overfitting the data. An example of this issue can be seen with the isolated elliptical galaxy NGC 720. This galaxy has a massive hot halo that has been observed once with Suzaku, twice with XMM-Newton, and several times with Chandra. One of the Chandra observations was studied by \citet{Humphrey2006}, who used the spectral method in eight annuli to measure the hot gas density profile out to 90 kpc; extrapolating their density profile to 300 kpc yields a hot halo mass of $1\times10^{11} M_{\odot}$, which implies that the galaxy is missing about 1/2 of its expected cosmological allotment of baryons. The rest of the data is analyzed in \citet{Humphrey2011}, with slight changes in their model, yielding a hot halo mass of $3\times10^{11} M_{\odot}$ within $R_{200}$, which implies that the galaxy is baryon complete. Based on the statistical errors quoted in these papers, this is about a $3\sigma$ discrepancy in the mass. About half the discrepancy is caused by the addition of the Suzaku data, and half by the changes in their modeling; however, at a smaller radius like $100$ kpc, the discrepancy is still nearly $3\sigma$, and this difference is caused almost entirely by their modeling. We will discuss this galaxy in much more detail in section 5. Finally, there are also situations where the spectral method and the spatial method yield different conclusions. A notable example is in the estimation of galaxy cluster density profiles near the virial radius - where the cluster emission is much fainter and the systematic uncertainties in these methods become more important. There has been some debate over the putative detection of a flattening in the radial decrease of the hot gas density profile, and this debate seems largely to fall along the lines between these two methods. The flattening was first observed with Suzaku, using spectral methods (\citealt{George2009}, \citealt{Simionescu2011}); spatial methods, using ROSAT, have often not confirmed this result (\citealt{Ettori2009}, \citealt{Eckert2011}, \citealt{Eckert2012}). Thus, while seems to be general consensus about a number of other features in the gas properties near the virial radius (such as decreasing temperature and flattening entropy), more work needs to be done to understand the behavior of the gas surface brightness and density (and therefore the total gas mass and baryon fraction of the cluster). Relatedly, discrepancies at the 5\%-15\% level for derived parameters such as luminosity, temperature, and pressure have also been noted by \citet{Rozo2012} in samples of clusters analyzed with different techniques (including combinations of spectral and spatial). In this paper we explore a potential improvement to the spatial method, taking into account both vignetted and unvignetted backgrounds based entirely on in-field data. This approach is similar to the use of the off field-of-view (OFOV) events for XMM-Newton, but can also be used for Chandra and Suzaku imaging, for which these events are not generally available. Adding an unvignetted component to the background model allows the entire image to be flat-fielded simultaneously, which is important for precise measurements of faint signals. The rest of the paper is devoted to taking the first steps towards full image modeling for X-ray astronomy. X-ray observations possess an unusual advantage over other wavelengths in that one records time, position, and energy measurements for each event. Most of this information is discarded when producing an image or a spectrum from an events file, but we argue that computational power has evolved to the point where it is no longer necessary to discard this information. We think the ultimate goal, which is outside the scope of this paper, is to study the X-ray events file directly instead of binning it spatially or spectrally to produce an image or a spectrum. With a good model for the events file that includes both energy and position (and potentially time), one could combine the best features of both spatial and spectral analysis, while minimizing the systematic errors associated with each. Here we implement a much more limited form of image modeling, using only spatial information in one energy band (typically 0.5-2.0 keV). We construct a simple model with both vignetted and unvignetted backgrounds. We then discuss the likelihood function for the image, and compare several different forms for the likelihood function before settling on one function. We show that this likelihood function is able to recover the input parameters in a variety of simulated images. Finally, we apply the method to the case of NGC 720 and show that even this limited form of image modeling offers significant advantages over traditional spectral fitting. Errors are quoted at $1\sigma$ unless otherwise noted.	In this paper, we have accomplished several goals. We argued that neither spectral fitting nor spatial fitting, at their current levels of sophistication, are adequate for the analysis of extended faint X-ray emission below about 1/5 of the background. We have also argued that an ideal approach would study X-ray observations at the level of the events file, and therefore make use of the spectral and spatial information simultaneously. In order to take the first steps towards such an approach, we presented an improvement to spatial fitting, wherein we model the entire image within a single energy band (0.5-2.0 keV). This discards more detailed energy information but can be extended in this direction in future work. We showed that a typical X-ray background can be decomposed into vignetted and unvignetted components. These components have different spectral shapes as well as different spatial distributions; in a given energy band, the different spatial distributions of these backgrounds can be used to constrain their relative contributions to an image. We introduced two methods of performing this decomposition. The {\it nonparametric} method excises a region around the source and fits the rest of the image in order to estimate the ratio of vignetted and unvignetted backgrounds, with no assumptions about the spatial distribution of the source emission other than that it must be localized to one region in the image which can be excluded. The nonparametric method will therefore not work for observations of diffuse emission which fills all or most of the field. The {\it parametric} method introduces a parametric model for the extended source (we examine $\beta$-models in this paper), and therefore is not limited to sources of small angular extent. We explored a number of possible likelihood functions for comparing the models to data. We decided upon a hybrid pixel-by-pixel likelihood function, with pixels binned where the psf is largest (and the correspond point source detectability the poorest). We showed that this likelihood function can reliably recover the shape of the background across the full spectrum from 100\% unvignetted to 100\% vignetted. This method works with either \verb"wavdetect" or with our manual point source detection algorithm, although we find the best results by masking point sources if either method detects them. We tested both methods for simulated extended sources observed with Chandra, in both the ACIS-I and ACIS-S configurations. We showed that we can recover the source emission well, recovering $\beta$ (a measure of the slope of the surface brightness profile) to 10\% accuracy and the total number of source counts to 25\% accuracy, for a source about half as bright as NGC 720. The method works better if the source is placed somewhat off-axis, so that its surface brightness profile can more easily be distinguished from the instrumental vignetting profile. Finally we applied our method to the isolated elliptical galaxy NGC 720. Out of two previous studies of the density profile of its hot halo, one study (Humphrey et al. 2006) seems to predict an X-ray surface brightness profile that matches reasonably well with the observed data, while the other study (Humphrey et al. 2011) systematically over-predicts the X-ray surface brightness. Both studies only measure density or surface brightness in a handful of annuli ($8-11$), however, so the inferred surface brightness profiles have large uncertainties - much larger than are warranted by examination of the surface brightness profiles directly. We argue this is an inherent feature of the spectral method, and is one reason why combined spatial and spectral analysis is preferable. The surface brightness profile predicted by our models offers a much better fit to the data. With this method, we are also able to trace the source emission to well below a tenth of the background - a significant improvement over previous methods. Both our parametric and non-parametric methods largely agree with each other, although our parametric model is insufficiently detailed to capture deviations from the single-component $\beta$-model which we observe at radii $\gapprox3$'. The implied mass of the hot halo extrapolated to 300 kpc is $\lapprox 1.5\times10^{11} M_{\odot}$. This is insufficient by about a factor of three to bring the galaxy to baryonic closure, in contrast to the conclusions of \citet{Humphrey2011}. We instead estimate that, after accounting for the stars and the hot gas, NGC 720 is missing over half of its baryons. This brings the galaxy into closer agreement with our estimates in \citet{Anderson2013} of the amount of hot gas around typical L* galaxies based on stacked images from the ROSAT All-Sky Survey.	14	3	1403.0584	Extended X-ray emission can be studied by spatial surface brightness measurements or by spectral analysis, but the two methods can disagree at low intensity levels. Here we present an improved method for spatial analysis that can be extended to include spectral information simultaneously. We construct a model for the entire image in a given energy band and generate a likelihood function to compare the model to the data. A critical goal is disentangling vignetted and unvignetted backgrounds through their different spatial distributions. Employing either maximum likelihood or Markov Chain Monte Carlo, we can derive probability distributions for the source and background parameters together, or we can fit and subtract the background, leaving the description of the source non-parametric. We calibrate this method against a variety of simulated images, and apply it to Chandra observations of the hot gaseous halo around the elliptical galaxy NGC 720. We follow the emission below a tenth of the background and infer a hot gas mass within 35 kpc of 4-5 × 10<SUP>9</SUP> M <SUB>⊙</SUB>, with some indication that the profile continues to at least 50 kpc and that it steepens. We derive stronger constraints on the surface brightness profile than previous studies that employed the spectral method, and we show that the density profiles inferred from these studies are in conflict with the observed surface brightness profile. Contrary to a previous claim, we find that the X-ray halo does not contain the full complement of missing baryons within the virial radius.	false	[ "spatial surface brightness measurements", "background parameters", "spectral analysis", "spectral information", "spatial analysis", "low intensity levels", "previous studies", "probability distributions", "the observed surface brightness profile", "Markov Chain Monte Carlo", "the surface brightness profile", "missing baryons", "NGC", "simulated images", "Extended X-ray emission", "Chandra observations", "×", "the density profiles", "the X-ray halo", "their different spatial distributions" ]	12.713084	8.106464	-1
437478	[ "Agúndez, Marcelino", "Parmentier, Vivien", "Venot, Olivia", "Hersant, Franck", "Selsis, Franck" ]	2014A&A...564A..73A	[ "Pseudo 2D chemical model of hot-Jupiter atmospheres: application to HD 209458b and HD 189733b" ]	119	[ "Univ. Bordeaux, LAB, UMR 5804, 33270, Floirac, France ; CNRS, LAB, UMR 5804, 33270, Floirac, France;", "Université de Nice-Sophia Antipolis, Observatoire de la Côte d'Azur, CNRS UMR 6202, BP 4229, 06304, Nice Cedex 4, France", "Instituut voor Sterrenkunde, Katholieke Universiteit Leuven, Celestijnenlaan 200D, 3001, Leuven, Belgium", "Univ. Bordeaux, LAB, UMR 5804, 33270, Floirac, France; CNRS, LAB, UMR 5804, 33270, Floirac, France", "Univ. Bordeaux, LAB, UMR 5804, 33270, Floirac, France; CNRS, LAB, UMR 5804, 33270, Floirac, France" ]	[ "2014ApJ...794..134W", "2014FaDi..168....9V", "2014PASA...31...43B", "2015A&A...577A..33V", "2015A&A...580A..12L", "2015ExA....40..329T", "2015ExA....40..469V", "2015ExA....40..481P", "2015Icar..257..163H", "2016A&A...587A.149V", "2016A&A...589A..52V", "2016A&A...594A..48L", "2016A&A...594A..69D", "2016A&A...595A..36A", "2016AJ....152..203L", "2016ApJ...817...17G", "2016ApJ...820...86M", "2016ApJ...821....9K", "2016ApJ...824..137Z", "2016ApJ...829...52F", "2016ApJ...829...66M", "2016ApJ...832...41M", "2016ApJS..224....9R", "2016MNRAS.460..855H", "2016MNRAS.461.1053M", "2016MNRAS.463..771M", "2016SSRv..205..285M", "2017A&A...600A..10M", "2017A&A...604A..79M", "2017AJ....153...68S", "2017ApJ...850..199W", "2017ApJ...851L..26D", "2017Icar..291....1C", "2017MNRAS.472..447M", "2018A&A...612A.105D", "2018A&A...614A.126L", "2018AJ....155...29W", "2018AJ....155..150M", "2018AJ....156...17K", "2018ApJ...853..138B", "2018ApJ...855L..31D", "2018ApJ...862...31T", "2018ApJ...863..183K", "2018ApJ...869...28D", "2018ApJ...869..107M", "2018ExA....46..101V", "2018ExA....46..135T", "2018MNRAS.474.5158G", "2018PASP..130k4402B", "2018arXiv180307163Z", "2018exha.book.....P", "2018fdpm.book..327M", "2018haex.bookE.116P", "2019A&A...623A.161C", "2019A&A...624A..58V", "2019AJ....158..217W", "2019ApJ...871...56M", "2019ApJ...872....1R", "2019ApJ...880...14S", "2019ApJ...886...39C", "2019LNP...955.....L", "2019MNRAS.487.2242H", "2020A&A...633A..86S", "2020A&A...636A..68D", "2020A&A...637A..59A", "2020A&A...639A...3C", "2020A&A...639A..48S", "2020A&A...641A.178H", "2020AJ....160..260C", "2020AJ....160..288F", "2020ApJ...890..176V", "2020ApJ...899...27Z", "2020ApJ...901..110P", "2020ApJ...905..131L", "2020MNRAS.493..106I", "2020RAA....20...99Z", "2020SSRv..216..139S", "2021A&A...653A..73S", "2021A&A...656A..90K", "2021AJ....161...18M", "2021ApJ...908..101R", "2021ApJ...913...73C", "2021MNRAS.501...78P", "2021MNRAS.502.6201B", "2021MNRAS.504.2783S", "2021MNRAS.505.4515R", "2021MNRAS.505.5603B", "2021MNRAS.506.3186H", "2021SSRv..217...43M", "2021arXiv210404824T", "2021exbi.book...18M", "2022A&A...664A..56S", "2022A&A...667A..15K", "2022AJ....163....8L", "2022ApJ...934...74M", "2022ExA....53..279M", "2022JQSRT.28908295B", "2022MNRAS.512.4877B", "2022PSJ.....3....1L", "2023A&A...672A.110L", "2023A&A...673A.125S", "2023ApJ...956..125O", "2023ApJ...957...20Z", "2023ApJ...958..143S", "2023IAUS..370..275K", "2023MNRAS.519.3129Z", "2023MNRAS.522.2525A", "2023MNRAS.525.5146Y", "2023Natur.614..653A", "2023RemS...15..635P", "2023arXiv231117075P", "2024A&A...682A.150K", "2024A&A...686A..24B", "2024AJ....167..195C", "2024MNRAS.tmp.1438Y", "2024NatAs.tmp....1B", "2024NatAs.tmp...86B", "2024PSJ.....5...58L", "2024arXiv240408759P" ]	[ "astronomy" ]	12	[ "planets and satellites: atmospheres", "planets and satellites: individual: HD 189733b", "planets and satellites: composition", "planets and satellites: individual: HD 209458b", "Astrophysics - Earth and Planetary Astrophysics" ]	[ "1969MNRAS.145..171T", "1973ppi..book.....B", "1977GeoNr..31...11I", "1982JQSRT..27..437H", "1989ApJ...336..495B", "1989ApJ...341..549B", "1997A&A...324..185B", "1997JQSRT..57..819K", "1999ppa..conf.....Y", "2001JQSRT..68..235B", "2002A&A...385..166S", "2002A&A...390..779B", "2004AdSpR..34..256T", "2005A&A...436..719I", "2005ApJ...629L..45C", "2005Natur.434..740D", "2006ApJ...644..560D", "2006ApJ...649.1048C", "2006ApJ...652..746F", "2007MNRAS.379..641C", "2007Natur.447..183K", "2007Natur.448..169T", "2008A&A...481L..83L", "2008A&A...492..585D", "2008ApJ...673..513D", "2008ApJ...673..526K", "2008ApJ...675..817C", "2008ApJ...678.1419F", "2008ApJ...678.1436B", "2008ApJ...678L.129C", "2008ApJ...681.1646R", "2008ApJ...682..559S", "2008ApJ...684.1427M", "2008ApJ...686.1341C", "2008Natur.452..329S", "2008Natur.456..767G", "2009A&A...505..891S", "2009ARA&A..47..481A", "2009ApJ...690L.114S", "2009ApJ...691..866K", "2009ApJ...699..478D", "2009ApJ...699..564S", "2009ApJ...701L..20Z", "2009ApJ...703..769K", "2009ApJ...704.1616S", "2009ApJ...707...24M", "2009JGRE..11411008E", "2009JQSRT.110..533R", "2010ARA&A..48..631S", "2010ApJ...708..498T", "2010ApJ...709.1396F", "2010ApJ...711..111M", "2010ApJ...714.1334R", "2010ApJ...717..496L", "2010ApJ...719..341B", "2010ApJ...719.1421P", "2010ApJ...720..344L", "2010ApJ...720.1569K", "2010ApJ...723.1436C", "2010ApJ...725..261M", "2010JQSRT.111.2139R", "2010MNRAS.408.1689S", "2010MNRAS.409..963B", "2010Natur.463..637S", "2010Natur.465.1049S", "2010RPPh...73a6901B", "2011A&A...532A...6S", "2011ApJ...726...82C", "2011ApJ...737...15M", "2011ApJ...738...32L", "2011ApJ...738...71S", "2011MNRAS.411.2199G", "2011MNRAS.413.2380H", "2011MNRAS.418.2669H", "2012A&A...546A..43V", "2012A&A...548A..73A", "2012A&A...548A.128D", "2012ApJ...744...35W", "2012ApJ...745...77K", "2012ApJ...745...78R", "2012ApJ...747L..20M", "2012ApJ...750...96R", "2012ApJ...751...59P", "2012ApJ...751...87D", "2012ApJ...751..117M", "2012ApJ...754...22K", "2012ApJ...757...46H", "2012ApJ...758...36M", "2012MNRAS.420..170L", "2012MNRAS.422..753G", "2013A&A...554A..82D", "2013A&A...558A..91P", "2013ApJ...762...24S", "2013ApJ...763...25M", "2013ApJ...764..103R", "2013ApJ...767...76K", "2013ApJ...774...95D", "2013ApJ...776..134P", "2013ApJ...776L..25D", "2013MNRAS.432.1980R", "2013MNRAS.436L..35B", "2014A&A...562A..51V", "2014ApJ...781...68A", "2014ApJ...783...70L", "2015A&A...574A..35P" ]	[ "10.1051/0004-6361/201322895", "10.48550/arXiv.1403.0121" ]	1403	1403.0121_arXiv.txt	Strongly irradiated by their close host star, hot Jupiters reside in extreme environments and represent a class of planets without analogue in our solar system. This type of exoplanets, the first to be discovered around main-sequence stars, remains the best available to study through observations and challenge a variety of models in the area of planetary science (see reviews on the subject by \cite{bar2010} 2010; \cite{sea2010} 2010; \cite{bur2011} 2011; \cite{sho2011} 2011). In recent years, multiwavelength observations of transiting hot Jupiters have allowed scientists to put constraints on the physical and chemical state of their atmospheres. Among these hot Jupiters, the best characterized are probably HD~209458b and HD~189733b, which belong to some of the brightest and closest transiting systems, and for which primary transit, secondary eclipse, and phase curve measurements have been used to probe, though often with controversial interpretations, various characteristics of their atmospheres, such as the thermal structure (\cite{dem2005} 2005, 2006; \cite{knu2008} 2008; \cite{cha2008} 2008), winds and day-night heat redistribution (\cite{knu2007} 2007, 2012; \cite{cow2007} 2007; \cite{sne2010} 2010), and mixing ratios of some of the main molecular constituents (\cite{tin2007} 2007; \cite{swa2008} 2008, 2009a, 2009b, 2010; \cite{gri2008} 2008; \cite{sin2009} 2009; \cite{des2009} 2009; \cite{bea2010} 2010; \cite{gib2011} 2011; \cite{wal2012} 2012; \cite{lee2012} 2012; \cite{rod2013} 2013; \cite{dek2013} 2013). The ability of observations at infrared and visible wavelengths to characterize the physical and chemical state of exoplanet atmospheres has motivated the development of various types of theoretical models. On the one hand, there are those aiming at investigating the physical structure of hot-Jupiter atmospheres, either one-dimensional radiative models (\cite{iro2005} 2005; \cite{for2008} 2008; \cite{par2014} 2014) or three-dimensional general circulation models (\cite{cho2008} 2008; \cite{sho2009} 2009; \cite{hen2011a} 2011a,b; \cite{dob2012} 2012; \cite{rau2013} 2013; \cite{par2013} 2013). These models have shown how fascinating the climates of hot Jupiters are, with atmospheric temperatures usually in excess of 1000 K, and helped to understand some global observed trends. Some hot Jupiters are found to display a strong thermal inversion in the dayside while others do not (e.g. \cite{for2008} 2008). Strong winds with velocities of a few km s$^{-1}$ develop and transport the heat from the dayside to the nightside, reducing the temperature contrast between the two hemispheres. The circulation pattern in these planets is characterized by an equatorial superrotating eastward jet. On the other hand, the chemical composition of hot Jupiters has been investigated by one-dimensional models, which currently account for thermochemical kinetics, vertical mixing, and photochemistry (\cite{zah2009} 2009; \cite{lin2010} 2010, 2011; \cite{mos2011} 2011, 2013; \cite{kop2012} 2012; \cite{ven2012} 2012, 2014; \cite{agu2014} 2014). These models have revealed the existence of three different chemical regimes in the vertical direction. A first one at the bottom of the atmosphere, where the high temperatures and pressures ensure a chemical equilibrium composition. A second one located above this, where the transport of material between deep regions and higher layers occurs faster than chemical kinetics so that abundances are quenched at the chemical equilibrium values of the quench level. And a third one located in the upper atmosphere, where the exposure to ultraviolet (UV) stellar radiation drives photochemistry. The exact boundaries between these three zones depend on the physical conditions of the atmosphere and on each species. In addition to the retrieval of average atmospheric quantities from observations, there is a growing interest in the physical and chemical differences that may exist between different longitudes and latitudes in hot-Jupiter atmospheres, and in the possibility of probing these gradients through observations. Indeed, important temperature contrasts between different planetary sides of hot Jupiters, noticeably between day and night sides, have been predicted (\cite{sho2002} 2002), observed for a dozen hot Jupiters (see \cite{knu2007} 2007 for the first one), qualitatively understood (\cite{cow2011} 2011; \cite{per2013} 2013), and confirmed by three-dimensional general circulation models (e.g. \cite{per2012} 2012). These temperature gradients, together with the fact that photochemistry switches on and off in the day and night sides, are at the origin of a potential chemical differentiation in the atmosphere along the horizontal dimension, especially along longitude. On the other hand, strong eastward jets with speeds of a few km s$^{-1}$ are believed to dominate the atmospheric circulation in the equatorial regions, as predicted by \cite{sho2002} (2002), theorized in \cite{shopol2011} (2011), potentially observed by \cite{sne2010} (2010), and confirmed by almost all general circulation models of hot Jupiters. These strong horizontal winds are an important potential source of homogenization of the chemical composition between locations with different temperatures and UV illumination. The existence of winds and horizontal gradients in the temperature and chemical composition of hot-Jupiter atmospheres has mainly been considered from a theoretical point of view, although some of these effects can be studied through phase curve observations (\cite{for2006} 2006; \cite{cow2008} 2008; \cite{maj2012} 2012; \cite{dew2012} 2012), monitoring of the transit ingress and egress (\cite{for2010} 2010), and Doppler shifts of spectral lines during the primary transit (\cite{sne2010} 2010; \cite{mil2012} 2012; \cite{sho2013} 2013). The existence of horizontal chemical gradients has been addressed in the frame of a series of one-dimensional models in the vertical direction at different longitudes (e.g. \cite{mos2011} 2011). An attempt to understand the interplay between circulation dynamics and chemistry was undertaken by \cite{coo2006} (2006), who coupled a three-dimensional general circulation model of HD 209458b to a simple chemical kinetics scheme dealing with the interconversion between CO and CH$_4$. These authors found that, even in the presence of strong temperature gradients, the mixing ratios of CO and CH$_4$ are homogenized throughout the planet's atmosphere in the 1 bar to 1 mbar pressure regime. In our team, we have recently adopted a different approach in which we coupled a robust chemical kinetics scheme to a simplified dynamical model of HD~209458b's atmosphere (\cite{agu2012} 2012). In this approach the atmosphere was assumed to rotate as a solid body, mimicking a uniform zonal wind, while vertical mixing and photochemistry were neglected. We found that the zonal wind acts as a powerful disequilibrium process that tends to homogenize the chemical composition, bringing molecular abundances at the limb and nightside regions close to chemical equilibrium values characteristic of the dayside. Here we present an improved model that simultaneously takes into account thermochemical kinetics, photochemistry, vertical mixing, and horizontal transport in the form of a uniform zonal wind. We apply our model to study the interplay between atmospheric dynamics and chemical processes, and the distribution of the main atmospheric constituents in the atmosphere of the hot Jupiters HD~209458b and HD~189733b.	We have developed a pseudo two-dimensional model of a planetary atmosphere that takes into account thermochemical kinetics, photochemistry, vertical mixing, and horizontal transport, and allows one to calculate the distribution with altitude and longitude of the main atmospheric constituents. Horizontal transport was modeled through a uniform zonal wind and thus the model is best suited for studying the atmosphere of planets whose circulation dynamics is dominated by an equatorial superrotating jet, as is expected to be the case of hot Jupiters. We therefore applied the model to study the atmospheres of the well-known exoplanets HD~209458b and HD~189733b. We used the temperature structure from GCM simulations and parameterized the turbulent mixing in the vertical direction using an eddy coefficient profile, which was calculated by following the behavior of passive tracers in GCM simulations, a method that results in substantially lower eddy values, by a factor of 10-100 in HD~209458b and of 10-1000 in HD~189733b, than previous estimates based on cruder methods. \emph{Molecular abundances homogenized with longitude to values typical of the hottest dayside regions. --} We found that the distribution of molecules in the atmospheres of HD~209458b and HD~189733b is quite complex because of the interplay of the various (photo)chemical and dynamical processes at work, which form, destroy, and transport molecules throughout the atmosphere. Much of the distribution of the atmospheric constituents is driven by the strong zonal wind, which reaches speeds of a few km s$^{-1}$, and the limited extent of vertical transport, with relatively low eddy diffusion coefficients below 10$^9$ cm$^2$ s$^{-1}$ around the 1 bar pressure level, resulting in an important homogenization of molecular abundances with longitude, in particular in the atmosphere of HD~189733b, which lacks a stratosphere and has quite low eddy diffusion coefficients. Moreover, molecular abundances are quenched horizontally to values typical of the hottest dayside regions, and therefore the composition of the cooler nightside regions is highly contaminated by that of warmer dayside regions. In hot Jupiters with a temperature inversion, such as HD~209458b, the longitudinal homogenization of molecular abundances is not as marked as in planets lacking a stratosphere, such as HD~189733b. In general, the cooler the planet, the stronger the homogenization of the chemical composition with longitude. Furthermore, in cooler planets such as hot Neptunes orbiting M dwarfs (e.g., GJ 436b) the temperature contrast between day and nightsides decreases because the cooling rate scales with the cube of temperature (e.g., \cite{lew2010} 2010), and therefore the composition is expected to be even more homogeneous with longitude than in warmer planets such as HD~209458b and HD~189733b. However, unlike hot Jupiters, hot Neptunes may have an atmospheric metallicity much higher than solar (\cite{lin2011} 2011; \cite{mos2013} 2013; \cite{agu2014} 2014; \cite{ven2014} 2014), which makes it interesting to investigate the extent of the spatial variation of molecular abundances in their atmospheres. \emph{Low methane content. --} A major consequence of our pseudo two-dimensional chemical model is that methane reaches quite low abundances in the atmospheres of HD~209458b and HD~189733b, lower than the values calculated by previous one-dimensional models. The main reason for the low CH$_4$ abundance is that most of the atmosphere is contaminated by the hottest dayside regions, where the chemical equilibrium abundance of CH$_4$ is the lowest. The calculated mixing ratio of CH$_4$ in the dayside of HD~209458b is significantly below the values inferred from observations, which points to some fundamental problem in either the chemical model or the observational side. If the strength of vertical transport is substantially higher than in our nominal model, the calculated abundance of some molecules such as CH$_4$ and NH$_3$ would experience significant enhancement, especially in HD~189733b, although a conflict with observations would still exist regarding CH$_4$ in the dayside of HD~209458b. \emph{Variability of planetary spectra driven by thermal, rather than chemical, gradients. --} An important consequence of the strong longitudinal homogenization of molecular abundances in the atmospheres of HD~209458b and HD~189733b is that the variability of the chemical composition has little effect on the way the emission spectrum is modified with phase and on the changes of the transmission spectrum from the transit ingress to the egress. Temperature variations and not chemical gradients are therefore at the origin of these types of variations in the planetary spectra. Only the longitudinal variation of the abundance of CO$_2$, of nearly one order of magnitude, in the atmosphere of HD~209458b, is predicted to induce variations in the planetary spectra around 4.3 and 15 $\mu$m. We note, however, that an inhomogenous distribution of clouds and/or hazes (none of them included in our model) may induce important variations in the emission spectra with phase and in the transmission spectra from one limb to the other. These variations are best characterized at short wavelengths. Indeed, there is evidence of the presence of hazes in the atmosphere of HD~189733b (\cite{lec2008} 2008; \cite{sin2009} 2009), and an inhomogeneous distribution of clouds has recently been inferred for the hot Jupiter Kepler 7b (\cite{dem2013} 2013). The main drawback of our pseudo two-dimensional chemical model is the oversimplification of atmospheric dynamics, which is probably adequate for equatorial regions, but not at high latitudes. Ideally, GCM simulations coupled to a robust chemical network would provide an even more realistic view of the distribution of molecules in the atmospheres of HD~209458b and HD~189733b, but such calculations are very challenging from a computational point of view. Telescope facilities planned for the near or more distant future, such as the James Webb Space Telescope, Spica, and EChO, will be able to test some of the predictions of our pseudo two-dimensional model, in particular the low abundance of methane in the two planets and the important longitudinal homogenization of the chemical composition.	14	3	1403.0121	The high temperature contrast between the day and night sides of hot-Jupiter atmospheres may result in strong variations of the chemical composition with longitude if the atmosphere were at chemical equilibrium. On the other hand, the vigorous dynamics predicted in these atmospheres, with a strong equatorial jet, would tend to supress such longitudinal variations. To address this subject we have developed a pseudo two-dimensional model of a planetary atmosphere, which takes into account thermochemical kinetics, photochemistry, vertical mixing, and horizontal transport, the latter being modeled as a uniform zonal wind. We have applied the model to the atmospheres of the hot Jupiters HD 209458b and HD 189733b. The adopted eddy diffusion coefficients were calculated by following the behavior of passive tracers in three-dimensional general circulation models, which results in much lower eddy values than in previous estimates. We find that the distribution of molecules with altitude and longitude in the atmospheres of these two hot Jupiters is complex because of the interplay of the various physical and chemical processes at work. Much of the distribution of molecules is driven by the strong zonal wind and the limited extent of vertical transport, resulting in an important homogenization of the chemical composition with longitude. The homogenization is more marked in planets lacking a thermal inversion such as HD 189733b than in planets with a strong stratosphere such as HD 209458b. In general, molecular abundances are quenched horizontally to values typical of the hottest dayside regions, and thus the composition in the cooler nightside regions is highly contaminated by that of warmer dayside regions. As a consequence, the abundance of methane remains low, even below the predictions of previous one-dimensional models, which probably is in conflict with the high CH<SUB>4</SUB> content inferred from observations of the dayside of HD 209458b. Another consequence of the important longitudinal homogenization of the abundances is that the variability of the chemical composition has little effect on the way the emission spectrum is modified with phase and on the changes in the transmission spectrum from the transit ingress to the egress. These variations in the spectra are mainly due to changes in the temperature, rather than in the composition, between the different sides of the planet.	false	[ "warmer dayside regions", "such longitudinal variations", "strong variations", "vertical transport", "HD", "chemical equilibrium", "horizontal transport", "vertical mixing", "planets", "the hottest dayside regions", "the hot Jupiters HD", "longitude", "previous estimates", "the strong zonal wind", "the cooler nightside regions", "little effect", "changes", "the chemical composition", "work", "thermochemical kinetics" ]	7.380523	14.712821	108
483006	[ "Montgomery, J. D.", "Clemens, D. P." ]	2014ApJ...786...41M	[ "Near-infrared Polarimetry of the Edge-on Galaxy NGC 891" ]	12	[ "Institute for Astrophysical Research, Boston University, 725 Commonwealth Avenue, Boston, MA 02215, USA", "Institute for Astrophysical Research, Boston University, 725 Commonwealth Avenue, Boston, MA 02215, USA" ]	[ "2013pss5.book..641B", "2017A&A...601A..92P", "2017ARA&A..55..111H", "2018ARA&A..56..673G", "2018ApJ...862...87S", "2018PASP..130e5002D", "2019PASA...36...29A", "2020AJ....160..167J", "2022MNRAS.515.5698M", "2023A&A...673A.112P", "2023AJ....165..223K", "2023ApJ...953..113L" ]	[ "astronomy" ]	3	[ "galaxies: individual: NGC 891", "galaxies: ISM", "galaxies: magnetic fields", "infrared: galaxies", "magnetic fields", "polarization", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1949ApJ...109..471H", "1949Sci...109..165H", "1949Sci...109..166H", "1970MmRAS..74..139M", "1974ApJ...194..249W", "1975ApJ...196..261S", "1982AJ.....87..695W", "1983asls.book.....B", "1987MNRAS.224..299S", "1988Natur.336..341R", "1989ApJ...346..728J", "1990IAUS..140..245S", "1991ApJ...382..100S", "1992ApJ...389..602J", "1992ApJ...400..238S", "1992MNRAS.257..309D", "1993MNRAS.264L...7S", "1994A&A...284..777G", "1994A&A...290..384D", "1996ARA&A..34..155B", "1996MNRAS.278..519S", "1996Natur.379...47B", "1996QJRAS..37..297S", "1997A&A...318..700B", "1997AJ....114.1393J", "1997AJ....114.1405W", "1997AJ....114.2463H", "1997ApJ...477L..25W", "1997ApJ...491..140S", "1998A&A...331..894X", "1998MNRAS.299..189S", "1999MNRAS.303..207S", "2000AJ....119..644H", "2005MNRAS.358..481K", "2006A&A...448..133K", "2007A&A...471L...1K", "2007ApJ...668..918B", "2007MNRAS.378..910L", "2007PASP..119.1385C", "2009MNRAS.395...97W", "2009RMxAC..36...25K", "2011ApJ...737..103S", "2011ApJS..195...18R", "2011MNRAS.412.2396F", "2012ApJ...761L..28P", "2012ApJS..200...19C", "2012ApJS..200...20C", "2012ApJS..200...21C", "2012MNRAS.423..197B", "2013AJ....145...74C" ]	[ "10.1088/0004-637X/786/1/41", "10.48550/arXiv.1403.3587" ]	1403	1403.3587_arXiv.txt	Viewed externally and \mbox{edge-on}, how would the magnetic (B) field threading the cold interstellar medium (ISM) of the Milky~Way appear? How is the ISM \mbox{B-field} generated and sustained, and what is its relationship to the radio synchrotron traced \mbox{B-field} in the hot plasma? NGC~891 is a bright, nearby \citep[$d = 9.1$~Mpc,][]{2011ApJS..195...18R}, \mbox{edge-on} \citep[$i = 89.8\arcdeg$,][]{1998A&A...331..894X, 2005MNRAS.358..481K}, Milky~Way analog ideally suited and situated for answering these questions. Polarization studies and other methods have revealed that \mbox{B-fields} permeate the ISM of the Milky Way and other galaxies \citep{1949ApJ...109..471H, 1949Sci...109..165H,1949Sci...109..166H, 1996Natur.379...47B, 1997A&A...318..700B, 2013AJ....145...74C}. Radio synchrotron emission is partially linearly polarized, with its sky projected orientation being perpendicular to the \mbox{B-field} in the hot ISM. Thus, synchrotron polarization has been used extensively to study the \mbox{B-field} associated with the hot plasma of external galaxies \citep{1996ARA&A..34..155B}. However, long radio wavelengths suffer Faraday rotation and depolarization, and accurately correcting for these is difficult \citep{1998MNRAS.299..189S, 1999MNRAS.303..207S}. Many \mbox{edge-on} galaxies studied in the radio show a characteristic \mbox{X-shaped} polarization morphology, where the \mbox{B-field} turns from disk-parallel to disk-perpendicular in the outer halo \citep{1994A&A...284..777G, 2006A&A...448..133K}. Radio studies of NGC~891 have revealed such an X-shape \citep{1991ApJ...382..100S, 2009RMxAC..36...25K}. Background optical and near-infrared (NIR) starlight probes \mbox{B-fields} in the dusty ISM. Radiative torques align aspherical dust grains with the long axis mostly perpendicular to the \mbox{B-field} \citep{2007MNRAS.378..910L}. This net alignment causes dichroic extinction, which linearly polarizes background starlight to exhibit an orientation parallel to the plane-of-sky \mbox{B-field} projection. Dichroic polarization percentage increases with optical depth, though not always linearly \citep{1989ApJ...346..728J, 1992ApJ...389..602J}. In addition to dichroic polarization, another source of polarization is light singly (and less so, multiply) scattered by dust grains. This light is polarized perpendicular to the plane of scattering and rarely relates to the \mbox{B-field} orientation. Therefore, it is important to distinguish between dichroism and scattering in NIR and optical studies of \mbox{B-field}s to avoid incorrect inferences regarding the \mbox{B-field} morphology and other properties. Although dust grain scattering cross-sections are inversely proportional to wavelength \citep{1983asls.book.....B}, the relationship between scattering polarization percentage and wavelength remains unclear. \citet{1992ApJ...400..238S} studied the reflection nebulae NGC~7023 and NGC~2023 and found no significant difference between $V$- and $H$-band polarization percentage. Thus, NIR wavelengths can probe deeper into the disk of NGC~891 than optical wavelengths, but scattering may still be an important component of the light and a possible contaminant in NIR \mbox{B-field} studies. Modern, high sensitivity, wide-field NIR polarimetry instruments are now able to probe \mbox{B-fields} across entire galaxies \citep{2007PASP..119.1385C, 2012ApJ...761L..28P}. The $H$-band polarization of the central 60~arcsec of NGC~891 was measured by \citet{1997AJ....114.1393J}, who found weak polarizations with generally disk-parallel orientation. NIR polarimetry of \emph{entire}, Milky~Way analog disks will provide context for detailed studies of the Milky~Way \mbox{B-field} \citep{1970MmRAS..74..139M, 2012ApJS..200...19C}. This current study sought to obtain $H$-band observations that are sufficiently sensitive to measure the polarization at arcsecond resolution out to nearly the full extent of the disk of NGC~891. These are the deepest NIR polarization observations of an entire \mbox{edge-on} galaxy to date. The observations revealed about 1\% polarization across the entire disk of NGC~891. This $P$ value is weaker than the \citet[][hereafter JKD]{1992ApJ...389..602J} model prediction for NGC~891. Polarization ``null points,'' where $P$ falls nearly to zero, perhaps related to the null points predicted by \citet{1997AJ....114.1405W} in their model of polarized galaxy light, were detected. However, the observed null points are not located symmetrically about the galaxy center, as was predicted. Therefore, the null points are not due to radiative transfer effects alone but may relate to galaxy structure such as spiral arms. There is strong polarization for positions well off the disk mid-plane, across the entire extent of NGC~891. This difference between mid-plane and off-plane polarizations might be due to disk starlight scattered by the extraplanar dust known to be present \citep{2000AJ....119..644H}. The new NIR polarimetric observations and the data reduction procedures are described in Section~\ref{sec:observations}. In Section~\ref{sec:analysis}, general polarization trends along the major and minor axes of NGC~891 are analyzed, polarization maps are presented, and radio and NIR P.A.s are compared. Section~\ref{sec:discussion} discusses how galaxy structure may influence polarization, compares these observations to polarization radiative transfer models, and compares the NIR and optical polarization properties.	\label{sec:discussion} The new $H$-band polarization data presented in this study reveal generally different P.A.s than seen for optical polarizations \citep{1996MNRAS.278..519S}. In addition, the properties of the NIR null point(s) contradict the \citet{1997ApJ...477L..25W} polarization model for an \mbox{edge-on} galaxy. What causes these differences? Several observations indicate the presence of extraplanar dust blown off the disk by supernovae and winds \citep{2000AJ....119..644H, 2007ApJ...668..918B, 2009MNRAS.395...97W}. Do the distributions of dust and starlight, or the presence of spiral arms, influence the observed polarizations? The new $H$-band polarization data revealed several north-south (N-S) asymmetries. These include the discovery of polarization null point(s) in the NE disk but not SW disk, strong polarization through the NE and central dust lane but weak polarization in the SW dust lane, disk-parallel P.A.s in the NE but slightly offset P.A.s in the south, and better agreement between NIR and radio P.A. values in the northern disk than in the southern disk. \subsection{Galaxy Structure} \subsubsection{Dust and Stellar Distributions} NGC~891 has a significant extraplanar dust component with an estimated vertical scale height of about 2~kpc \citep{2007A.A...471L...1K, 2009MNRAS.395...97W}, while the stellar distribution scale height is 0.35~kpc \citep{1998A&A...331..894X}. Much of this extraplanar dust has visual extinctions of order unity \citep[$\tau_{H} \approx 0.15$]{2000AJ....119..644H}. The minor axis polarization percentage profile (Figure~\ref{fig:minorPro}.a) shows a strong increase in polarization strength beginning about 0.35~kpc off the dusty midplane. The relative strength and vertical height at which this polarization trend begins indicate a likely origin in disk starlight scattered by the extraplanar dust and into the line-of-sight. Thus, toward the disk midplane, where there is significant background starlight that can be polarized by dichroism, NIR polarizations likely do trace the \mbox{B-field}. However, far off the mid-plane, where there is little background starlight, polarizations are likely dominated by scattering and do not trace the \mbox{B-field}. The regions where scattering likely dominates the polarization are marked in Figure~\ref{fig:minorPro} by gray shading. \subsubsection{Spiral Arms} The N-S asymmetries noted above may be related to spiral arms. Many other N-S asymmetries were previously observed in NGC~891. \citet{1998A&A...331..894X} and \citet{2007A.A...471L...1K} observed asymmetries in the NIR brightness and H$\alpha$ light, respectively, and both concluded these indicated the presence of a spiral arm on the near side (relative to the center of NGC~891) of the NE disk. However, similar asymmetries in the \ion{H}{1} emission \citep{1997ApJ...491..140S} and at other radio wavelengths \citep{1994A&A...290..384D} suggest that spiral structure and extraplanar dust might not be responsible for the observed asymmetries. If spiral arms influence the observed distribution of light even in \mbox{edge-on} galaxies, where they cannot be directly observed, then the polarization asymmetries---especially the null points--- may be similarly related to the influence of spiral arms. Section~\ref{sec:spiralNulls} further discusses the possible connection between spiral arms and polarization null points. \subsection{Dynamo and Polarization Models} \subsubsection{$\alpha\Omega$ Dynamo} The new NIR polarization data reported here confirm a significant difference between the galaxy major axis P.A. and the $H$-band polarization P.A., as seen toward the galaxy central region by \citet{1997AJ....114.1393J}, across most of the full extent of NGC~891. This P.A. offset is at odds with the predictions of typical $\alpha\Omega$~dynamo models \citep{1988Natur.336..341R}. Also, the P.A. offset cannot be due to polarization effects within the Milky~Way, as the $H$-band foreground extinction to NGC~891 is only 0.029~mag \citep{2011ApJ...737..103S}. Interestingly, \citet{2012ApJS..200...21C} found that the median Galactic Position Angle (GPA) for background starlight polarizations in the Galactic plane is 75\arcdeg, a 15\arcdeg\ offset from the expected disk-parallel GPA of 90\arcdeg. Thus, it is possible that significant departures from toroidal B-fields are not uncommon. \subsubsection{Polarization and Radiative Transfer} \citet{1997ApJ...477L..25W} modeled the expected polarization due to scattering and dichroism for a galaxy with purely toroidal \mbox{B-fields} for several galaxy inclination angles. In his model, light from the bright galaxy center was scattered by dust in the disk to produce disk-perpendicular polarizations throughout the galaxy \citep[Figure~1]{1997ApJ...477L..25W} while light from the stellar disk passed through magnetically aligned dust grains to produce disk-parallel polarizations. In the emergent polarization images of the model \mbox{edge-on} galaxies, a symmetric pair of polarization null points appeared in the model disk at the projected offsets from the galaxy center where the orthogonal scattering and dichroic polarizations canceled. The null points were associated with a transition from disk-parallel polarization in the central region to disk-perpendicular polarization further from the bulge. The data presented in Figure~\ref{fig:majorPro} of the present study show at least one highly significant NIR polarization null point, yet the polarization orientation revealed in Figure~\ref{fig:regA} is disk-parallel on \emph{both} sides of that null point. This result argues against any transition from dichroic to scattering polarization across this null point. The \citet{1997ApJ...477L..25W} model did not include spiral structure in the distribution of starlight or dust or in the magnetic field geometry. This simple model predicted the occurrence of a \emph{symmetric pair} of null points in the disk, one on either side of the galaxy center. However, Figure~\ref{fig:majorPro} shows that no null points are observed in the SW disk. Rather, there is a continuous strip of low polarization along the SW disk dust lane. Thus, the observed null point(s) are different in nature than those predicted by the \citet{1997ApJ...477L..25W} model. \citet{1997AJ....114.1405W} specifically modeled the galaxies NGC~4565 and NGC~891 at $V$-band and $H$-band wavelengths, using the \citet{1997ApJ...477L..25W} formalism, and they compared the model results to existing polarization observations in these bands. They found that observed polarizations of NGC~891 in both $V$- and $H$-band could not be replicated by their model. They attributed the discrepancies to strong winds dragging the \mbox{B-field} up into the halo, as evidenced by the detection of extraplanar dust blown off the disk \citep{1997AJ....114.2463H}. \subsubsection{Weak Polarization Percentage} NGC~891 shows particularly weak $P$ across its entire disk. The JKD model predicts $\sim 2\%$ polarization based on the $E(H-K)$ color of the dust lane of NGC~891. \citet{1997AJ....114.1393J} noted that NGC~4565, another normal \mbox{edge-on} galaxy, shows polarizations much stronger than NGC~891 and in agreement with the JKD model. He suggested two possible explanations. The low $P$ might be due to a crossed-polaroid effect from toroidal B-fields in the inner disk and poloidal B-fields in the outer disk. Section~\ref{sec:opticalVsNIR} discusses why this is unlikely. The other \citet{1997AJ....114.1393J} explanation requires tangled B-fields on scales smaller than 100~pc but not on any larger scales. This B-field configuration seems untenable given Figures~4, 5, and 6 here. A third explanation relates to scattered bulge light. Scattering of bulge light by dust in the disk mid-plane is not the dominant source of polarization along the major axis, although some scattered light ought to be present. This scattered light would be polarized with disk-perpendicular P.A. and could provide a crossed-polaroid effect, which would diminish the net polarization percentage. However, the bulge to disk luminosity ratio of NGC~891 is not particularly high, so it is unclear why this would be the case in this galaxy but not in others. The low $P$ for NGC~891 is yet to be successfully explained. \subsection{Polarization Null Points and Spiral Arms} \label{sec:spiralNulls} If the polarization null points found here are not caused by the mechanism described by \citet{1997ApJ...477L..25W}, then what \emph{does} cause null points, and what influences their locations? Could the null points be related to spiral arms? \citet{2011MNRAS.412.2396F} observed that the radio synchrotron traced \mbox{B-field} is oriented parallel to the spiral arms of M51 (a \mbox{face-on} galaxy). If the \mbox{B-field} in the \mbox{edge-on} NGC~891 is similarly directed along its spiral arms, then for lines-of-sight with significant optical depth inside of a spiral arm, the net plane-of-sky component of the \mbox{B-field} will be small, and the resulting dichroic polarization will be weak. Assuming a smooth exponential distribution of dust in the disk, with a conservative $V$-band optical depth $\tau_{V} = 3$ through the galaxy pole, model $H$-band optical depths were estimated. With a nearly \mbox{edge-on} galaxy inclination, unit optical depth at $H$-band corresponds to a path length of $\sim$6~kpc through the disk. Thus, depending on the scale lengths of the dust distribution, $H$-band observations may probe deeply enough to reach points where the ``magnetic spiral arms'' are tangent to the line-of-sight, producing polarization nulls having disk-parallel polarization signatures flanking each null. \subsection{Optical vs. NIR polarization} \label{sec:opticalVsNIR} \citet{1996MNRAS.278..519S} interpreted their observed $V$-band, predominantly disk-perpendicular, polarizations as evidence for poloidal \mbox{B-fields} in the outer disk of NGC~891. However, the scattering only, \mbox{edge-on} galaxy in the \citet{1997ApJ...477L..25W} models produced disk-perpendicular polarizations (even in the central region, where \citeauthor{1996MNRAS.278..519S} argued scattering could not be responsible for the observed polarization P.A.). Furthermore, if the \citeauthor{1996MNRAS.278..519S} disk-perpendicular polarizations were caused by poloidal \mbox{B-fields} in the outer disk, then the $H$-band polarizations presented here for the extreme NE and SW disk should also have been disk-perpendicular, but they are not. Instead, bulge light scattered by the disk dust into the line-of-sight might be the dominant source of optical wavelength polarized light, in which case optical polarization is unlikely to be tracing the \mbox{B-field} of NGC~891. \citet{2012ApJ...761L..28P} found no $H$-band polarization from the \mbox{face-on} galaxy M51. If the lack of polarization was due to insufficient NIR optical depth through the galaxy to produce detectable polarization, then there may be similarly insufficient dichroic optical depth in the visible bands, assuming a normal Serkowski law relationship \citep{1975ApJ...196..261S, 1982AJ.....87..695W} between polarization and wavelength. If so, the centro-symmetric optical polarizations observed in \mbox{face-on} galaxies \citep{1987MNRAS.224..299S, 1992MNRAS.257..309D, 1993MNRAS.264L...7S, 1996QJRAS..37..297S} may also be caused by scattering and not dichroism. This would leave radio as the sole means for probing the \mbox{B-field} of \mbox{face-on} galaxies and weaken the reliability of optical polarizations for probing the \mbox{B-fields} of even \mbox{edge-on} galaxies.	14	3	1403.3587	The edge-on galaxy NGC 891 was probed using near-infrared (NIR) imaging polarimetry in the H band (1.6 μm) with the Mimir instrument on the 1.8 m Perkins Telescope. Polarization was detected with a signal-to-noise ratio greater than three out to a surface brightness of 18.8 mag arcsec<SUP>-2</SUP>. The unweighted average and dispersion in polarization percentage (P) across the full disk were 0.7% and 0.3%, respectively, and the same quantities for polarization position angle (P.A.) were 12° and 19°, respectively. At least one polarization null point, where P falls nearly to zero, was detected in the northeast disk but not the southwest disk. Several other asymmetries in P between the northern and southern disk were found and may be related to spiral structure. Profiles of P and P.A. along the minor axis of NGC 891 suggest a transition from magnetic (B) field tracing dichroic polarization near the disk mid-plane to scattering dominated polarization off the disk mid-plane. A comparison between NIR P.A. and radio (3.6 cm) synchrotron polarization P.A. values revealed similar B-field orientations in the central-northeast region, which suggests that the hot plasma and cold, star-forming interstellar medium may share a common B-field. Disk-perpendicular polarizations previously seen at optical wavelengths are likely caused by scattered light from the bright galaxy center and are unlikely to be tracing poloidal B-fields in the outer disk.	false	[ "dominated polarization", "synchrotron polarization P.A. values", "dichroic polarization", "polarization position angle", "Polarization", "polarization percentage", "Disk-perpendicular polarizations", "NIR P.A.", "spiral structure", "P.A.", "similar B-field orientations", "plane", "poloidal B-fields", "mid", "the outer disk", "-", "the full disk", "the northeast disk", "scattered light", "optical wavelengths" ]	12.923348	9.336575	-1
437532	[ "Breslau, Andreas", "Steinhausen, Manuel", "Vincke, Kirsten", "Pfalzner, Susanne" ]	2014A&A...565A.130B	[ "Sizes of protoplanetary discs after star-disc encounters" ]	62	[ "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany" ]	[ "2015A&A...577A.115V", "2015ApJ...806L..10D", "2015MNRAS.447..577M", "2015MNRAS.449.1996D", "2015PASP..127..961A", "2015PhyS...90f8001P", "2016A&A...594A..30W", "2016A&A...594A..53B", "2016ApJ...828...48V", "2016MNRAS.455.3086X", "2016MNRAS.457..313P", "2016MNRAS.457.3593F", "2016MNRAS.457.4218J", "2016MNRAS.460L.109M", "2017A&A...599A..91B", "2017A&A...602A..12R", "2017A&A...602A..52W", "2017A&A...604A..91W", "2017MNRAS.471.2334X", "2018A&A...610A..33P", "2018A&A...611A..88L", "2018ApJ...863...45P", "2018ApJ...868....1V", "2018ApJ...869L..44K", "2018ComAC...5....2V", "2018MNRAS.475.2314W", "2018MNRAS.477..325B", "2018MNRAS.478.2700W", "2018exha.book.....P", "2019A&A...621A.101B", "2019A&A...622A..69P", "2019A&A...624A.110F", "2019A&A...625A..59J", "2019ARep...63..190M", "2019ComAC...6....3B", "2019MNRAS.483.4114C", "2019MNRAS.485.1489W", "2019MNRAS.488.1366L", "2019MNRAS.490.5678C", "2020AJ....159..106W", "2020ARA&A..58..483A", "2020AcA....70...53J", "2020ApJ...897...60P", "2020ApJ...898..140H", "2020MNRAS.491..504C", "2020MNRAS.491..903W", "2020RSOS....701271P", "2021A&A...651A..38P", "2021A&A...652A.144P", "2021A&A...653A.168J", "2021ApJ...915..131L", "2021ApJ...921...90P", "2021ApJ...923..221O", "2021ApJS..257...19H", "2021RNAAS...5...10J", "2022MNRAS.512..861N", "2022MNRAS.515.2837W", "2022NatAs...6..837L", "2023EPJP..138...11C", "2023MNRAS.518.5620W", "2023MNRAS.520.5331W", "2023MNRAS.520.6159C" ]	[ "astronomy" ]	8	[ "protoplanetary disks", "planets and satellites: formation", "galaxies: star clusters: general", "Astrophysics - Astrophysics of Galaxies" ]	[ "1981ARA&A..19..137P", "1993MNRAS.261..190C", "1994ApJ...424..292O", "1995ApJ...455..252H", "1996MNRAS.278..303H", "1997MNRAS.287..148H", "1998AJ....115..263O", "1998ApJ...499..758J", "2000AJ....119.2919B", "2001ApJ...553L.153H", "2001Icar..153..416K", "2001MNRAS.325..449S", "2003AJ....126.1916P", "2003ARA&A..41...57L", "2003ApJ...592..986P", "2005A&A...437..967P", "2005A&A...441..195V", "2005ApJ...629..526P", "2006ApJ...641..504A", "2006ApJ...642.1140O", "2006ApJ...653..437M", "2006Icar..184...59B", "2007A&A...462..193P", "2007ApJ...659..705A", "2009ApJ...700.1502A", "2010ARA&A..48...47A", "2011A&A...532A.120L", "2011ARA&A..49...67W", "2011MNRAS.411..859M", "2011MNRAS.418.1272J", "2012A&A...538A..10S", "2012A&A...546L...1D", "2013A&A...549A..82P", "2013ApJ...769..150C", "2013ApJ...771..129A" ]	[ "10.1051/0004-6361/201323043", "10.48550/arXiv.1403.8099" ]	1403	1403.8099_arXiv.txt	\label{sec:intro} Stars form through the gravitational collapse of dense cores within molecular clouds. Due to the angular momentum conservation during the collapse, most young stars are initially surrounded by a disc consisting of gas and dust. These protoplanetary discs are usually treated as unperturbed by their surrounding. However, as most stars do not form in isolation but in star cluster environments ({{Lada} \& {Lada}} 2003; {{Porras} {et~al.}} 2003) discs are prone to be affected by external processes like photoevaporation from massive stars (e.g. {{Johnstone} {et~al.}} 1998; {{Scally} \& {Clarke}} 2001) or gravitational interactions with other cluster members (e.g. {{Moeckel} \& {Bally}} 2006; {{Craig} \& {Krumholz}} 2013). In this paper we only concentrate on the effect of the latter. The consequences of gravitational interactions - or encounters - on disc parameters, like the disc's mass, angular momentum and energy, have been intensively investigated analytically and numerically in the past (e.g. {{Clarke} \& {Pringle}} 1993; {Ostriker} 1994; {Heller} 1995; {{Hall} {et~al.}} 1996; {{Kobayashi} \& {Ida}} 2001; {Pfalzner} 2003; {{Olczak} {et~al.}} 2006; {{Pfalzner} \& {Olczak}} 2007{a}; {{Lestrade} {et~al.}} 2011; {{Steinhausen} {et~al.}} 2012). By contrast, the encounter induced alteration of the size of a protoplanetary disc has so far been investigated in less detail. The reason is that only for the order of a few hundred objects, disc sizes have so far been determined (for an overview see e.g. {{Williams} \& {Cieza}} 2011 and references therein). Deriving correlations between the disc size and the stellar environment seemed so far out of reach. However, with the advent of ALMA in the near future the sample size is likely to increase considerably and the influence of the environment on the disc size becomes testable. In the past, theoretical investigations have shown that encounters mainly affect the outer disc material. {{Clarke} \& {Pringle}} (1993) investigated the effect of parabolic encounters of equal-mass stars with a periastron distance of $1.25$ times the initial radius of the disc, $r_{\mathrm{init}} $, and varying inclinations. They found that in the most destructive encounter (prograde, coplanar) disc material down to $\approx 0.5$ times the periastron distance, $r_{\mathrm{peri}} $, becomes unbound. By varying the disc-size to periastron-distance ratio, {{Hall} {et~al.}} (1996) found, that in parabolic, prograde, coplanar encounters of equal mass stars the disc inside of $0.3\,r_{\mathrm{peri}} $ is almost unperturbed. {{Kobayashi} \& {Ida}} (2001) focused on the encounter induced increase of eccentricity and inclination of the disc material and the size of the largely unperturbed inner part of the disc. They found that many particles become unbound outside of $\approx 1/3$ of periastron distance after a prograde encounter of equal-mass stars on a parabolic orbit. Almost all previous investigations focused on a small parameter space, which was mostly restricted to encounters between equal-mass stars. Nevertheless, the above results are often generalised as encounters truncating discs to $1/2\text{--}1/3$ of the periastron distance (e.g. {{Brasser} {et~al.}} 2006; {{Adams} {et~al.}} 2006; {Adams} 2010; {{Jim{\'e}nez-Torres} {et~al.}} 2011; {{Malmberg} {et~al.}} 2011; {Pfalzner} 2013). There the dependence of the disc truncation on the mass of the perturbing star and other geometrical properties of the perturber orbit are disregarded. One typical context where the disc size after an encounter plays a key role is the formation of the solar system. The size of the solar system, with a density drop at $\approx 30$\,AU, is attributed to an encounter of the young solar system with another star at a periastron distance of $\approx 100$\,AU. This distance is used to determine the typical cluster environment in which the solar system might have developed. In addition, the disc size is a crucial parameter for investigations of the typical types of planetary systems in a given cluster environment. Therefore, {{de Juan Ovelar} {et~al.}} (2012) tried to estimate the sizes of discs after encounters considering both, the periastron distance and the mass ratio of the stars. They suggested two different approaches. The first one is based on the assumption that disc material is removed in an encounter at least to the point of gravitational force equilibrium between the stars at the time of periastron passage. In the other approach they suggested a transformation of the results for the disc-mass loss in encounters by {{Olczak} {et~al.}} (2006) into a reduction of the disc size. Both approaches assume that disc size change and the removal of disc material are strongly correlated. However, already {Hall} (1997) showed that during an encounter disc material can be moved inward due to loss of angular momentum. That way, the disc size can be reduced even when no mass is lost. More recently, it was demonstrated that a $3\text{--}5\,\%$ loss of disc angular momentum is common in ONC-like star clusters and this holds not only for the central high-density regions but even at the outskirts of clusters ({{Pfalzner} \& {Olczak}} 2007{b}). Therefore it is very likely, that changes of the disc size are a more common effect in clusters than the removal of disc material. Using an extended data base of star-disc encounters we show here that mass loss based approaches are not really suitable to determine the resulting disc size. After a short description of our numerical method in Sect.~\ref{sec:method}, a disc-size definition is given, which is comparable to observational size determinations. In Sect.~\ref{sec:results} we present the numerical results as well as a simple fit formula. We compare our results to previous work in Sect.~\ref{sec:comparison} and give a brief summary and conclusion in Sect.~\ref{sec:summary}.	\label{sec:discussion} Some approximations have been made in the model described above. First, we treated the discs by pure N-body methods while neglecting viscous forces. Viscous forces only play a role in the central areas of the discs. For typical viscosity values the effected area is the region within $ \lesssim 0.2\,r_{\mathrm{init}} $. Only in the most violent interactions the disc size is reduced to such low values. For these cases viscous forces would have to be in principle included. As can be seen in Table~\ref{tab:sizes}, this is, for example, the case for equal-mass stars in penetrating encounters closer than $ 0.7\,r_{\mathrm{init}} $. For typical disc sizes on the order of some $100$\,AU (e.g. {{Bally} {et~al.}} 2000; {{Andrews} {et~al.}} 2009), such encounters are relatively rare in ONC-Type clusters ({{Olczak} {et~al.}} 2006). In addition, for such destructive encounters the material that remains bound is usually $<20\,\%$ of its initial mass and its structure does not resemble a disc as such. In these cases the definition of a disc size is anyway highly questionable. An additional approximation is that the case where only one of the stars was surrounded by a disc was considered. This has been done for simplicity of description. In principle, for disc-disc encounters the sizes could be larger because disc material could be transferred from the perturbing star to the disc hosting star and replenish the disc. {{Pfalzner} {et~al.}} (2005{a}) showed that in case both stars are surrounded by discs, mostly an additive approach can be used as usually captured material is deposited close to the star. In addition, the amount of captured mass is very small compared to the disc mass as such. Therefore in most cases the final disc size is little influenced by captured material. The exception is again the case where the encounter is very violent and the remnant disc mass very low. Furthermore, the outcome of a disc-disc encounter could be influenced by the dependence of the disc mass and size on the stellar mass. Observations give no conclusive answer to this question. If the disc size scales in some way with the stellar mass, the situation can occur that the perturber has a much higher mass, and therefore has a bigger and possibly more massive disc (see e.g. {{Andrews} {et~al.}} 2013). In this case the amount of material captured by the primary could be comparable to or even higher than that of its remaining disc. We have chosen the point of highest surface density gradient as definition of the disc size (see Sect.~\ref{sec:method_size}). However, after an encounter the region where the disc's surface density drops significantly spans some range (see Fig.~\ref{fig:disc_size}). The inner boundary of this region would correspond better to the disc sizes obtained by SED fits, where the size is usually defined as the point, where the surface density transits from a power-law to something steeper. If we had defined our disc sizes as the inner boundary of the density-drop region, the results would be smaller than with our actual definition. Conversely, if defining the size as the point, where the surface density falls below a certain threshold (compare Sect.~\ref{sec:method_size}), the sizes would be somewhat bigger than our results. However, results obtained with each of these definitions would usually not differ from our results by more than $\approx 0.1\,r_{\mathrm{init}} $. Our definition is therefore a robust mean value between several possible disc size definitions for perturbed discs. Since the problem of star disc encounters scales with the periastron ratio, one would expect that the results can be normalised to this ratio $r_{\mathrm{peri}} /r_{\mathrm{init}} $. This is indeed the case, but for more intuitive comparison with observations we have chosen the absolute presentation of the results. The actual values in Table~\ref{tab:sizes} in the Online Material are normalised to the periastron ratio. The differences between the simulation results and the fit formula increase for decreasing mass ratios, becoming significant ($\gtrsim 5\%$) for ${m_{\mathrm{12}}} \lesssim 5$. This is because in the close encounters needed to change the disc size for low mass ratios the redistribution of the disc material is non-linear. Therefore the fit formula with linear dependence on the periastron distance can describe this effect only approximately. It is also important to note, that in contrast to close, penetrating encounters, the disc size change in distant encounters is not mainly caused by a real truncation of the disc. Mostly, the size change is dominated by the redistribution of disc material towards the host star. This redistribution results in smaller disc sizes even when no material is lost at all. \begin{figure} \centering \includegraphics[width=\hsize]{parameter_space_overview} \caption{The parameter space of encounters occurring in a star cluster like the ONC (full height and open to the right) compared to the parameter space of our simulations (grey area) and the validity area of our fit formula (grey, not ruled area). For further details see text.} \label{parameter_space} \end{figure} Figure~\ref{parameter_space} shows where our fit formula can be applied. The axes are chosen to cover the encounter parameter space typical for an ONC-like star cluster. The grey area shows the parameter space covered by this study. Right from the solid bent line, encounters have almost no effect on the mass, angular momentum or size of a disc. Right from the dashed bent line, it can be assumed that encounters have also a negligible effect. The horizontal, dashed line depicts the parameter space of {{Hall} {et~al.}} (1996) while the black dot marks the simulations by {{Kobayashi} \& {Ida}} (2001). In the black shaded area the remaining discs have sizes below $0.2\,r_{\mathrm{init}} $ or have less than $10\,\%$ of their initial particles. Since the simulations in this area may be influenced by low resolution or the hole in our discs, we excluded them. The remaining grey area depicts the parameter range for which our fit formula is valid. {{Kobayashi} \& {Ida}} (2001) developed an analytical estimate for the size of the inner disc region, where the velocity dispersion of the disc material after an encounter is still low enough to form planets (see their Eq.~(31)). They obtained \begin{align} r_{\mathrm{planet}} \propto \left( \frac{{m_{\mathrm{12}}} +1}{{m_{\mathrm{12}}} ^2} \right)^{1/4} \left(\frac{r_{\mathrm{peri}} }{r_{\mathrm{init}} } \right)^{5/4}. \end{align} Even though the size of this planet forming region is defined differently than our disc size, the mass dependence is similar to the ${m_{\mathrm{12}}} ^{-0.32}$ dependence of Eq.~\eqref{eq:fitformula}. However, the periastron dependence of their analytical approximation is stronger than that of our numerical results. \label{sec:summary} In the dense stellar environments of young clusters, tidal interactions can change the sizes of protoplanetary discs. Performing N-body simulations of such encounters, we confirm earlier results by {{Kobayashi} \& {Ida}} (2001) that close ($\le 100$\,AU) encounters between equal-mass solar-type stars lead to the shrinking of initially $100$\,AU-sized discs to $30\text{--}50$\,AU. Naturally this represents only a special case, in clusters the wide spectrum of encounter partners and periastra has to be taken into account. In this paper we investigated the disc-size change by encounters for the entire parameter space spanned by mass ratio and periastron distance typically covered in clusters. The central result of our extensive numerical parameter study is that the disc size after a prograde, parabolic encounter is a simple function of the periastron distance $r_{\mathrm{peri}} $ and the mass ratio ${m_{\mathrm{12}}} $ of the two stars of the form \begin{align} r_{\mathrm{final}} = 0.28 \cdot r_{\mathrm{peri}} \cdot {m_{\mathrm{12}}} ^{-0.32}. \notag \end{align} Prograde, parabolic encounters are the most destructive type of encounters. Inclined and/or hyperbolic encounters lead to less mass loss and therefore larger disc sizes. However, the parameter dependencies in these types of encounters span a wide parameter range. We will investigate this extended parameter range in a follow up study (in preparation). The disc sizes after a star-disc encounter as obtained with the here presented disc-size definition would not necessarily correspond to the final size of a potentially developing planetary system. The reasons are that on long time scales (several Myr) viscosity leads to an increase in disc size. Simultaneously highly eccentric particles probably become recircularised through viscous processes as they pass the inner parts of the disc while being close to their periastron. Previous work on disc sizes after encounters was often motivated by the search for the solar birth environment. Here the sudden density drop in the mass distribution at $30\text{--}50$\,AU is interpreted as the result of a close fly-by of another star during the formation phase of the solar system. Considering only encounters between equal mass stars an encounter distance of $100\text{--}150$\,AU was deduced. However, recent results show that encounters of the early solar system with less or more massive stars are at least just as likely as with another solar-mass star. Our results show now, that any parabolic, prograde encounter which fulfils the relation \begin{align} 0.28 \cdot r_{\mathrm{peri}} \cdot {m_{\mathrm{12}}} ^{-0.32} = 30\text{--}50~\mathrm{AU} \end{align} can lead to a solar-system size disc. So far the dependencies of protoplanetary disc sizes on parameters like stellar mass etc. are observationally not well constrained. With the advent of ALMA this will quickly change. Thus it will be also possible to determine whether dense stellar environments have a significant influence on the disc sizes and the forming planetary systems. The here derived dependencies will be a useful tool to determine the corresponding encounter events. \bibpunct{(}{)}{;}{a}{}{,} %	14	3	1403.8099	Most stars do not form in isolation, but as part of a star cluster or association. These young stars are initially surrounded by protoplanetary discs. In these cluster environments tidal interactions with other cluster members can alter the disc properties. Besides the disc frequency, its mass, angular momentum, and energy, the disc's size is particularly prone to being changed by a passing star. So far the change in disc size has only been investigated for a small number of very specific encounters. Several studies investigated the effect of the cluster environment on the sizes of planetary systems like our own solar system, based on a generalisation of information from this limited sample. We performed numerical simulations covering the wide parameter space typical of young star clusters, to test the validity of this approach. Here the sizes of discs after encounters are presented, based on a size definition that is comparable to the one used in observational studies. We find that, except for encounters between equal-mass stars, the usually applied estimates are insufficient. They tend to severely overestimate the remaining disc size. We show that the disc size after an encounter can be described by a relatively simple dependence on the periastron distance and the mass ratio of the encounter partners. This knowledge allows us, for example, to pin down the types of encounter possibly responsible for the structure of today's solar system. <P />Appendix A is available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201323043/olm">http://www.aanda.org</A>	false	[ "disc size", "young star clusters", "protoplanetary discs", "encounter", "encounters", "other cluster members", "discs", "Most stars", "planetary systems", "observational studies", "Several studies", "the remaining disc size", "the disc size", "a star cluster", "our own solar system", "todays solar system", "information", "equal-mass stars", "tidal interactions", "the encounter partners" ]	9.33705	13.2685	-1
734946	[ "Postnov, Konstantin A.", "Yungelson, Lev R." ]	2014LRR....17....3P	[ "The Evolution of Compact Binary Star Systems" ]	305	[ "Sternberg Astronomical Institute, Moscow M.V. Lomonosov State University, 13 Universitetskij Pr., 119992, Moscow, Russia", "Institute of Astronomy of the Russian Academy of Sciences, 48 Pyatnitskaya St., 119017, Moscow, Russia" ]	[ "2014ApJ...795L...9B", "2014MNRAS.445.3239P", "2014PhRvD..90h4043M", "2015AIPC.1693c0004C", "2015ARep...59...25I", "2015ApJ...798L..19M", "2015ApJ...803...41M", "2015ApJ...804...21G", "2015ApJ...810...58S", "2015AstL...41..114K", "2015AstL...41..797M", "2015MNRAS.448..541J", "2015MNRAS.450L..85M", "2015MNRAS.452.1970W", "2015MNRAS.452.2897P", "2015MNRAS.453.3024C", "2015MNRAS.454.4411D", "2015PhRvD..92j2001K", "2015PhRvD..92l4007D", "2015PhRvL.115w1102F", "2016ARNPS..66...23F", "2016ApJ...818L..22A", "2016ApJ...828...77V", "2016ApJ...830L..18B", "2016ComAC...3....6T", "2016GReGr..48...95M", "2016JPhCS.675c2020M", "2016JPhCS.706e2016C", "2016MNRAS.455...17V", "2016MNRAS.456.4089B", "2016MNRAS.457.1636B", "2016MNRAS.458.2634M", "2016MNRAS.460L...1S", "2016MNRAS.461.2747K", "2016MNRAS.461.4329I", 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"2013LRR....16....7G", "2013MNRAS.429.1602Y", "2013MNRAS.429.2143C", "2013MNRAS.429.2361L", "2013MNRAS.429L.104Z", "2013MNRAS.430..274F", "2013MNRAS.430..996L", "2013MNRAS.430.1746G", "2013MNRAS.430.2281N", "2013MNRAS.431..372C", "2013MNRAS.431.1812T", "2013MNRAS.431.2778K", "2013MNRAS.432.1264K", "2013MNRAS.432.1640W", "2013MNRAS.432.2048K", "2013MNRAS.432.2378P", "2013MNRAS.433.1114Y", "2013MNRAS.433.2884I", "2013MNRAS.433L..20M", "2013MNRAS.434..102S", "2013MNRAS.434.1355J", "2013MNRAS.435..187N", "2013MNRAS.436.3380E", "2013Natur.500..547T", "2013PASA...30...46C", "2013PhRvD..87h4006F", "2013PhRvD..87j4028G", "2013PhRvD..87l3007B", "2013PhRvD..88d3011P", "2013PhyU...56..714G", "2013PrPNP..68....1R", "2013RvMP...85..245B", "2013Sci...339..433I", "2013Sci...340..170W", "2013arXiv1309.6635K", "2013pss4.book..397H", "2013sepa.book.....I", "2013sepp.book.....I", "2014A&A...562A..95G", "2014A&A...563A..16N", "2014A&A...563A..83C", "2014A&A...564A.134M", "2014AJ....147..129S", "2014ARA&A..52..107M", "2014ARA&A..52..487S", "2014ARep...58..113P", "2014ASPC..490..287N", "2014ApJ...781..104G", "2014ApJ...782...27K", "2014ApJ...783L...8G", "2014ApJ...784..119R", "2014ApJ...785...28K", "2014ApJ...785...61S", "2014ApJ...785..105M", "2014ApJ...786L..11M", "2014ApJ...788...75R", "2014ApJ...789L...5K", "2014ApJ...792..123B", "2014LRR....17....2B", "2014MNRAS.437..649S", "2014MNRAS.437.1681M", "2014MNRAS.437.2894C", "2014MNRAS.437L..66S", "2014MNRAS.438...14D", "2014MNRAS.438.3399B", "2014MNRAS.438L..26K", "2014MNRAS.439..757G", "2014MNRAS.439.2765S", "2014MNRAS.439.2848C", "2014MNRAS.439.3064M", "2014MNRAS.440..504S", "2014MNRAS.440.1274C", "2014MNRAS.440.1498S", "2014MNRAS.441..532D", "2014MNRAS.441.1186D", "2014MNRAS.442.1079J", "2014MNRAS.443.1849S", "2014MNRAS.445.1912C", "2014Natur.505..378C", "2014PhR...539...49K", "2014PhRvD..89f4056M", "2015A&A...583A.140S", "2015HiA....16...51V", "2015MNRAS.448..928K", "2015MNRAS.448.2362P", "2015MNRAS.454L..61D", "2018LRR....21....3A" ]	[ "10.12942/lrr-2014-3", "10.48550/arXiv.1403.4754" ]	1403	1403.4754_arXiv.txt			14	3	1403.4754	We review the formation and evolution of compact binary stars consisting of white dwarfs (WDs), neutron stars (NSs), and black holes (BHs). Mergings of compact-star binaries are expected to be the most important sources for forthcoming gravitational-wave (GW) astronomy. In the first part of the review, we discuss observational manifestations of close binaries with NS and/or BH components and their merger rate, crucial points in the formation and evolution of compact stars in binary systems, including the treatment of the natal kicks, which NSs and BHs acquire during the core collapse of massive stars and the common envelope phase of binary evolution, which are most relevant to the merging rates of NS-NS, NS-BH and BH-BH binaries. The second part of the review is devoted mainly to the formation and evolution of binary WDs and their observational manifestations, including their role as progenitors of cosmologically-important thermonuclear SN Ia. We also consider AM CVn-stars, which are thought to be the best verification binary GW sources for future low-frequency GW space interferometers.	false	[ "compact binary stars", "binary GW sources", "BH", "BHs", "binary evolution", "compact stars", "massive stars", "neutron stars", "binary systems", "binary WDs", "GW", "close binaries", "future low-frequency GW space interferometers", "evolution", "NS-BH and BH-BH binaries", "crucial points", "NS and/or BH components", "compact-star binaries", "NS-BH", "observational manifestations" ]	7.414688	2.814745	39
485359	[ "Chakrabarty, Deepto", "Tomsick, John A.", "Grefenstette, Brian W.", "Psaltis, Dimitrios", "Bachetti, Matteo", "Barret, Didier", "Boggs, Steven E.", "Christensen, Finn E.", "Craig, William W.", "Fürst, Felix", "Hailey, Charles J.", "Harrison, Fiona A.", "Kaspi, Victoria M.", "Miller, Jon M.", "Nowak, Michael A.", "Rana, Vikram", "Stern, Daniel", "Wik, Daniel R.", "Wilms, Jörn", "Zhang, William W." ]	2014ApJ...797...92C	[ "A Hard X-Ray Power-law Spectral Cutoff in Centaurus X-4" ]	57	[ "MIT Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Space Sciences Laboratory, University of California, Berkeley, CA 94720, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Department of Astronomy, University of Arizona, Tucson, AZ 85721, USA", "Observatoire Midi-Pyrénées, Université de Toulouse III - Paul Sabatier, F-31400 Toulouse, France; CNRS, Institut de Recherche en Astrophysique et Planetologie, F-31028 Toulouse, France", "Observatoire Midi-Pyrénées, Université de Toulouse III - Paul Sabatier, F-31400 Toulouse, France; CNRS, Institut de Recherche en Astrophysique et Planetologie, F-31028 Toulouse, France", "Space Sciences Laboratory, University of California, Berkeley, CA 94720, USA", "Division of Astrophysics, National Space Institute, Technical University of Denmark, DK-2800 Lyngby, Denmark", "Space Sciences Laboratory, University of California, Berkeley, CA 94720, USA; Lawrence Livermore National Laboratory, Livermore, CA 94550, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Department of Physics, McGill University, Montreal, PQ H3A 2T8, Canada", "Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA", "MIT Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA 91109, USA", "Astrophysics Science Division, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Dr. Karl-Remeis-Sternwarte and Erlangen Centre for Astroparticle Physics, Universität Erlangen-Nürnberg, D-96049 Bamberg, Germany", "Astrophysics Science Division, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA" ]	[ "2014ApJ...791...77T", "2014MNRAS.442..372K", "2015ApJ...801...10A", "2015ApJ...806..148B", "2015ApJ...807...33P", "2015ApJ...807...52A", "2015ApJ...809...13D", "2015MNRAS.449.2803D", "2015MNRAS.451.2071D", "2015MNRAS.452.3475B", "2015MNRAS.454.1371W", "2015arXiv150103534T", "2016ApJ...821..103R", "2016ApJ...825..132H", "2016ApJ...826..149B", "2016ApJ...826..162E", "2016ApJ...831..184B", "2016MNRAS.456.4001W", "2017ApJ...845....8L", "2017ApJ...850..155S", "2017JApA...38...49W", "2017MNRAS.464..398D", "2017MNRAS.465L..10D", "2017MNRAS.466.4074P", "2017MNRAS.471.3494Z", "2017MNRAS.472.2742P", "2017PASJ...69...23O", "2017SSRv..212..429C", "2017arXiv171203949V", "2018ASSL..457..185D", "2018ApJ...853..157H", "2018ApJ...854...58A", "2018ApJS..235...26Z", "2018MNRAS.473.3789A", "2018MNRAS.475.2027V", "2018MNRAS.476..421S", "2018MNRAS.476.2230P", "2018MNRAS.477.2494V", "2018MNRAS.479.2777R", "2018MNRAS.479.3634M", "2018arXiv180602833D", "2019AN....340..226L", "2019MNRAS.483.4560Z", "2019MNRAS.488.4477D", "2019cxro.book....6N", "2020A&A...638L...2P", "2020MNRAS.492..615Q", "2020MNRAS.496.2704Q", "2021ASSL..461....1M", "2021MNRAS.501.1453P", "2021MNRAS.502.3870Q", "2022ApJ...930...20B", "2022ApJ...933..240G", "2022ApJ...934..142S", "2022MNRAS.515.3838M", "2022MNRAS.516.2641V", "2024MNRAS.531.1653V" ]	[ "astronomy" ]	12	[ "accretion", "accretion disks", "binaries: close", "stars: individual: Cen X-4", "stars: neutron", "X-rays: binaries", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1969ApJ...157L.157C", "1970ApJ...159L..57E", "1972ApJ...171L..87B", "1975A&A....39..185I", "1979IAUC.3369....1H", "1979rpa..book.....R", "1980ApJ...236L..55C", "1980ApJ...238..287J", "1980ApJ...240L.137M", "1980ApJ...241..779K", "1984ApJ...278..270B", "1984ApJ...283..694K", "1984ApJ...283..710K", "1988ApJ...335L..75H", "1988Natur.335..801K", "1989A&A...210..114C", "1990ARA&A..28..215D", "1993ApJ...402..593S", "1993ApJ...403..249A", "1994ApJ...423L..47S", "1994ApJ...434..570T", "1995ApJ...439..849Z", "1995ApJ...449..188H", "1995ApJ...450..876T", "1995ApJ...452..710N", "1996ASPC..101...17A", "1996ApJ...464L.127K", "1996ApJ...464L.139V", "1996ApJ...465..487V", "1997ApJ...490..605M", "1998A&ARv...8..279C", "1998ApJ...499L..65C", "1998ApJ...504L..95B", "1998PASJ...50..611A", "1998tbha.conf..148N", "1999ApJ...514..945R", "1999ApJ...520..276M", "1999ApJ...521..332P", "1999MNRAS.303L...1B", "2000ApJ...531..956M", "2000ApJ...541..908B", "2000ApJ...542..914W", "2001A&A...365L...1J", "2001A&A...365L..18S", "2001A&A...365L..27T", "2001A&A...377..955D", "2001A&A...378L...5Z", "2001ApJ...557..304M", "2001ApJ...560L..71B", "2001NewAR..45..449L", "2002ASPC..271...71A", "2002ApJ...577..346R", "2002MNRAS.334..233T", "2002apa..book.....F", "2003A&A...404L..43B", "2003ApJ...597..474C", "2003MNRAS.345.1271V", "2003Sci...299.1372S", "2004ApJ...601..474C", "2004ApJ...614L..49C", "2004MNRAS.354..666J", "2005A&A...440..775K", "2005A&A...444..905D", "2005AJ....130..759T", "2005ApJ...620..922C", "2006ARA&A..44...17G", "2006ApJ...644.1090H", "2006ApJ...646..304U", "2008MNRAS.391.1619D", "2009A&A...503..889K", "2009ApJ...703.2017W", "2009Sci...324.1411A", "2010A&A...515A..25D", "2010ApJ...720.1325C", "2010ApJ...722...88A", "2010MNRAS.406.1208D", "2011AAS...21714425P", "2011ApJ...737..103S", "2011ApJ...742...97B", "2011MNRAS.410.1886S", "2011MNRAS.414.3006C", "2012MNRAS.420..416D", "2013A&A...550A..89D", "2013ApJ...770..103H", "2013ApJ...772....7G", "2013IAUS..291..127R", "2013MNRAS.432.2366W", "2013MNRAS.433.1362C", "2013MNRAS.436.2465B", "2013Natur.501..517P", "2014AN....335..313R", "2014ATel.5890....1R", "2014ApJ...781L...3P", "2014ApJ...789...40B", "2014ApJ...790...39S", "2014ApJ...791...77T", "2014ApJ...792...48W", "2014MNRAS.438..251L", "2014MNRAS.438.2634C", "2014MNRAS.440..504S", "2014MNRAS.441...86L", "2014MNRAS.441.1825B", "2014PASJ...66...10S" ]	[ "10.1088/0004-637X/797/2/92", "10.48550/arXiv.1403.6751" ]	1403	1403.6751_arXiv.txt	Low-mass X-ray binaries (LMXBs) consist of a neutron star (NS) or black hole (BH) accreting from a low-mass ($\lesssim~1~M_{\odot}$) stellar companion via Roche-lobe overflow. They may be divided into two categories: persistent accretors with X-ray luminosity $L_{\rm x}\gtrsim$10$^{36}$~erg~s$^{-1}$, and transient systems. Transient LMXBs undergo recurrent bright ($L_{\rm x}\gtrsim$10$^{36}$~erg~s$^{-1}$) outbursts lasting days to weeks and then return to long intervals of X-ray quiescence ($L_{\rm x}\lesssim$10$^{34}$~erg~s$^{-1}$) lasting months to years. The long-term average mass accretion rate of the transients is thus substantially lower than in the persistent systems, owing to their low duty cycle. Transient behavior is understood to arise from a thermal instability in the outer accretion disk wherein the viscosity (and thus the mass accretion rate $\dot M$ through the disk) jumps to a higher value when a critical surface density is reached as the disk fills up \citep[see][and references therein]{las01}. The persistent LMXBs avoid this instability because their higher accretion rates lead to increased X-ray heating, keeping the disks permanently ionized \citep{van96,kkb96}. For the NS systems, the 0.5--10 keV X-ray spectrum of quiescent LMXB transients typically consists of two components: a low-energy (``soft'') $\sim$0.1~keV thermal component, and a high-energy (``hard'') power-law component with photon index $1<\Gamma<2$, where photon flux $dN/dE\propto E^{-\Gamma}$. The soft component is generally well fit by a hydrogen atmosphere model for the NS. The leading explanation for the energy source of the soft component is a deep crustal heating model \citep{bbr98} in which the emission is powered by heat injected into the NS crust by pycnonuclear reactions driven by accretion during transient outbursts. In this model, the contribution of quiescent accretion is negligible. X-ray spectroscopy of soft thermal emission in quiescent NS transients has been used to infer NS radii \citep{bbr98,rbb+99,gsw+13} and to study the thermal relaxation of NS crusts \citep[see][and references therein]{wdp13}. However, a possible problem for such studies is that accretion may not have completely shut off during quiescence, as suggested by the detection of quiescent variability in the two brightest quiescent NS/LMXBs, Aql X-1 \citep{rbb+02} and Cen X-4 \citep{cis+04,bcb+13}. There has been considerable debate as to whether this variability is primarily in the soft thermal component, the hard power-law component, or both \citep[e.g.,][]{rbb+02,cs03,cwh+05}. The origin of the hard power-law tail is unclear. It is not predicted by the deep crustal heating model \citep{bbr98}. Two explanations have been discussed: synchrotron shock emission from a radio pulsar wind and Comptonization of the soft thermal photons by a hot corona \citep{ccm+98}. The synchrotron model is of particular interest given the recent confirmation that NS/LMXBs can turn on as radio pulsars at low accretion rates \citep{asr+09,pfb+13}. A difficulty in discriminating between different models has been the absence of knowledge about how high the power-law component extends in energy, owing to lack of sufficient observational sensitivity above 10~keV. The recent launch of the {\em NuSTAR} hard X-ray telescope provides the first opportunity to explore this question. The ideal target with which to address this is Cen X-4 = X1455$-$314 (Galactic coordinates $l=332.2^\circ$, $b=23.9^\circ$), the brightest quiescent NS/LMXB. It was discovered in 1969 in the 3--12~keV band with the {\em Vela 5A/5B} satellites during an extremely bright ($\sim$20 Crab at peak) X-ray outburst lasting over two months \citep{ceb69,ebc70}. A second bright ($\sim$4 Crab at peak) X-ray outburst was detected in 1979 \citep{khs80} along with counterparts in the optical \citep{cmg80} and radio \citep{hje79,hch+88}, but the source has been in X-ray quiescence ($\sim 10^5 \times$ fainter) ever since. Bright X-ray flashes, now understood as thermonuclear (type I) X-ray bursts, were observed around the time of both the 1969 and 1979 outbursts \citep{bce72,mik+80}, conclusively establishing the source as a NS and setting an upper limit on the distance of 1.2$\pm$0.3~kpc \citep{civ+89}. A third burst may have been observed by the {\em Apollo 15} lunar mission in 1971 \citep{kil09}. The presence of thermonuclear bursts indicates that the surface dipole magnetic field is weak, with $B_{\rm surf}\lesssim 10^{10}$~G \citep{jl80} and most likely $\sim 10^8$~G (by analogy with other type I bursters). Optical photometry and spectroscopy indicate that the binary companion V822~Cen is a K3-7 V dwarf, the binary period is 15.1~hr, and the binary mass ratio is $q=0.1755$ \citep{civ+89,tcm+02,dcc+05,swd14}. The best-fit NS mass is $1.94^{+0.37}_{-0.85} M_\odot$ \citep{swd14}. Given the proximity and high Galactic latitude of the source, it has extremely low interstellar extinction and absorption, allowing more sensitive observations in the ultraviolet and soft X-ray bands than usually possible for LMXBs \citep{brd+84,mr00,pg11,cbd+13}. The integrated values along this line of sight are $A_V=0.362$ \citep{sf11} and $N_{\rm H}\approx 9\times 10^{20}$~cm$^{-2}$ \citep{dl90,kbh+05}. Cen X-4 has been observed extensively during X-ray quiescence since the 1979 outburst, with deep X-ray spectra in the 0.5--10~keV range previously obtained on six occasions since 1994 \citep[see][and references therein; see also \S5.2]{cbm+10}. A long-term daily monitoring campaign with {\em Swift} recently demonstrated that the thermal and power-law components vary together on time scales from days to months, with no spectral change observed and each component contributing roughly half the flux \citep{bcb+13}. These authors concluded that the quiescent X-ray emission in Cen X-4 is primarily generated by accretion. In this paper, we present the first sensitive hard X-ray observation of Cen X-4 in quiescence with {\em NuSTAR}, obtained simultaneously with a deep {\em XMM-Newton} soft X-ray observation. We describe the observations in \S2 and our spectral analysis and results in \S3. We interpret our results in \S4 and discuss their implications in \S5.	\subsection{The spectral cutoff in Cen X-4} We have measured a cutoff in the hard X-ray power-law spectrum of Cen X-4 which can be fit with an exponential cutoff at around 10~keV or a bremsstrahlung spectrum with $kT_e=18$~keV. This is the first detection of a power-law cutoff in this source class, and it finally permits a more detailed investigation of the origin of the hard component in quiescent NS/LMXBs. We were able to rule out both thermal Comptonization and synchrotron shock emission as the origin of the spectral cutoff. Instead, the hard X-ray spectrum can be understood as bremsstrahlung emission, arising either from a hot, optically thin corona above the NS atmosphere or from hot electrons in an optically thin RIAF. The NS atmosphere scenario has the advantage that it can easily explain why the soft and hard emission varies together on short time scales, while the RIAF scenario has the advantage that it can self-consistently account for the observed luminosity. The 18~keV electron temperature for the {\tt bremss} model is much lower than either the $\sim$50~keV electron temperature predicted for the hot layer above a NS atmosphere \citep{dds01} or the $\gtrsim$100~keV electron temperature expected in a RIAF flow around a black hole \citep{mq97}. This may be due to Compton cooling of the bremsstrahlung electrons by the soft X-ray photons from the NS atmosphere, in which case we would expect $T_e$ to depend upon the soft X-ray luminosity $L_{\rm soft}$. The absence of a detectable Compton emission component is not problematic. The Compton radiative power density is \begin{equation} P_C = \sigma_T\,n_e\,\left(\frac{4 kT_e}{m_e c^2}\right) \left(\frac{L_{\rm soft}}{4\pi r^2}\right) . \label{eq:compt_radpow} \end{equation} The resulting inverse Compton luminosity is \begin{eqnarray} L_{C, {\rm atm}} & = & 5\times 10^{29}\,h_3^{1/2}\, \left(\frac{kT}{\mbox{\rm 18 keV}}\right) \nonumber \\ & & \quad \times \left(\frac{L_{\rm soft}/L_E}{2\times 10^{-6}}\right) \mbox{\rm\ erg~s$^{-1}$} \end{eqnarray} for cooling above the NS atmosphere, or \begin{eqnarray} L_{C, {\rm RIAF}} & = & 3\times 10^{30}\, \left(\frac{\eta}{0.1}\right)^{-1} \left(\frac{kT}{\mbox{\rm 18 keV}}\right) \left(\frac{\dot M_{\rm NS}/\dot M_E}{4\times 10^{-6}}\right) \nonumber \\ & & \quad \times \left(\frac{L_{\rm soft}/L_E}{2\times 10^{-6}}\right) \mbox{\rm\ erg~s$^{-1}$} \end{eqnarray} for cooling in the accretion flow. In either case, the Compton luminosity is no more than a few percent of the overall source luminosity, and thus essentially undetectable in our spectrum. It is interesting to consider whether the flares observed in the X-ray light curve might be expected to affect the electron temperature (and thus the cutoff energy). The Compton cooling time scale is $t_C =(3/2)n_e kT_e/P_C$. For emission above the NS atmosphere, this yields \begin{equation} t_{C,{\rm atm}} = 1\times 10^{-2}\, M_{1.9}^{-1}\, \left(\frac{L_{\rm soft}/L_E}{2\times 10^{-6}}\right)^{-1} \mbox{\rm\ s} \,, \end{equation} so the cooling is nearly instantaneous. For this scenario, we might expect to measure spectral changes in the {\em NuSTAR} band ($\gtrsim$10~keV) during the flares, although we did not have sufficient signal-to-noise to do this with our observation. By contrast, emission in the RIAF gives \begin{equation} t_{C,{\rm RIAF}} = 1\times 10^{6}\, M_{1.9}^{-1}\, \left(\frac{r}{10^4 R}\right)^2 \left(\frac{L_{\rm soft}/L_E}{2\times 10^{-6}}\right)^{-1} \mbox{\rm\ s} \,, \end{equation} so that in this case short-term flaring behavior will not result in significant Compton cooling of the bremsstrahlung electrons in the RIAF flow. This might provide an avenue for discriminating between the two scenarios. On longer time scales (months to years), the fact that the slope of the hard X-ray power-law spectrum in Cen X-4 was observed to vary from epoch to epoch over the course of two decades \citep{cbm+10} makes it unlikely that the 18~keV bremsstrahlung spectrum we have measured is a constant feature of the source in quiescence. In fact, we can demonstrate that the hard X-ray cutoff energy in Cen X-4 is likely variable by noting that the departure of the hard X-ray spectrum from an unbroken power-law above $\simeq$6~keV is evident in our 2013 {\em XMM-Newton} spectrum alone, even without including the {\em NuSTAR} data. The shape of an 18~keV bremsstrahlung spectrum will show noticeable curvature below 10~keV. On the other hand, all previous deep observations of Cen X-4 in the 0.5--10~keV band are consistent with an unbroken hard X-ray power-law spectrum, indicating a higher cutoff energy for those observations. All of these observations occurred at significantly lower luminosity (see Figure~\ref{fig_gamma}). We would expect a lower thermal luminosity to result in reduced Compton cooling of the bremsstrahlung electrons and hence a higher electron temperature, consistent with a higher cutoff energy. \begin{deluxetable}{llcccc} \tablecaption{HARD X-RAY POWER-LAW SPECTRA OF DEEP CEN X-4 OBSERVATIONS \tablenotemark{a}\label{tbl-3}} \tablewidth{\textwidth} \tablehead{ & & \colhead{Exposure} & \colhead{$L_{\rm th}$\tablenotemark{b}}& & \\ \colhead{Start date} & \colhead{Mission} & \colhead{(ks)} & \colhead{($10^{32}$ erg s$^{-1}$)} & \colhead{$\Gamma$} & \colhead{Ref.} } \startdata 1994 Feb 27 & {\em ASCA} & 39 & 1.19(11) & 1.24(17) & 1 \\ 2001 Aug 20 & {\em XMM} & 53 & 1.50(5) & 1.41(5) & 1 \\ 2003 Mar 1 & {\em XMM} & 78 & 1.07(2) & 1.26(8) & 1 \\ 2009 Jan 16 & {\em Suzaku} & 147 & 0.29(2) & 1.69(17) & 1 \\ 2010 Aug 25 & {\em XMM} & 21 & 0.63(6) & 1.77(21) & 2 \\ 2010 Sep 4 & {\em XMM} & 23 & 0.67(1) & 1.62(10) & 2 \\ 2011 Jan 24 & {\em XMM} & 15 & 0.97(2) & 1.38(10) & 2 \\ 2011 Jan 31 & {\em XMM} & 14 & 0.31(1) & 1.94(19) & 2 \\ 2013 Jan 20 & {\em XMM}+{\em NuSTAR}\tablenotemark{c} & 27/114 & 3.8(1) & 1.56(5) & 3 \enddata \tablenotetext{a}{1$\sigma$ uncertainties in last digits shown in parentheses.} \tablenotetext{b}{0.5--10 keV thermal luminosity assuming $D$=1 kpc.} \tablenotetext{c}{Fit only to 0.3--10 keV data using {\tt tbabs(nsatmos+powerlaw)} model, with no power-law break.} \tablerefs{(1) Cackett et al. 2010 and references therein; (2) Cackett et al. 2013; (3) This work.} \tablecomments{All archival data fit to {\tt phabs(nsatmos+powerlaw)} model. {\em Chandra} observations excluded owing to possible photon pileup.} \end{deluxetable} \begin{figure}[t] \begin{center} \includegraphics[width=0.47\textwidth]{fig_gamma.eps} \end{center} \caption{Hard X-ray power-law photon index (2--10 keV) versus unabsorbed 0.5--10~keV thermal luminosity for deep observations of Cen X-4 (see Table~\ref{tbl-3}). The {\em XMM}/{\em NuSTAR} point is from a {\tt tbabs*(nsatmos+powerlaw)} fit to the 0.3--10~keV data only. The thermal luminosity $L_{\rm th}$ is computed for a distance of 1~kpc. For $L_{\rm th}\gtrsim 10^{32}\,D_{\rm kpc}^2$ erg~s$^{-1}$, the data are roughly consistent with a trend of steeper power-law slope for higher $L_{\rm th}$ as expected for bremsstrahlung emission, if we assume that $kT_e$ is reduced from 50--100~keV via Compton cooling by thermal photons. In energy bandpasses well below $kT_e$, a bremsstrahlung spectrum is a $\Gamma=1$ power law. The large $\Gamma$ measured at the lowest thermal luminosities may indicate a transition to synchrotron shock emission at extremely low $\dot M$. \label{fig_gamma}} \end{figure} Moreover, we would expect the 2--10~keV power-law slope to vary systematically with the thermal luminosity $L_{\rm th}$. For energy bandpasses far below $kT_e$, a bremsstrahlung spectrum is a $\Gamma=1$ power-law \citep[see, e.g.][]{rl79}; as one approaches the cutoff at $kT_e$ from below, the effective $\Gamma$ over a fixed bandpass increases as the spectrum begins to gradually roll over. Thus, if we assume that $kT_e\gtrsim$~50--100~keV in the absence of Compton cooling, then a Compton-cooled bremsstrahlung model predicts that $\Gamma$ should be close to 1 at low $L_{\rm th}$ and should increase as $L_{\rm th}$ rises and $kT_e$ falls. This relationship will eventually break down when $kT_e$ gets sufficiently low, because an unbroken power-law no longer provides even a rough fit to a sharp spectral cutoff. The archival data roughly support this picture for $L_{\rm th}\gtrsim 10^{32}\,D_{\rm kpc}^2$ erg~s$^{-1}$. In Figure~\ref{fig_gamma}, we plot $\Gamma$ versus $L_{\rm th}$ for deep quiescent observations of Cen X-4 made since 1994, including our observation\footnote{We exclude observations made with {\em Chandra}/ACIS, which may be subject to pileup effects. We note that the $\Gamma$ values found by \citet{cbm+10} for these observations are much smaller than for any other observations, as would be expected if there is significant pileup. We will reexamine these data elsewhere.}. These observations are listed in Table~\ref{tbl-3}; the archival spectra are collected from \citet{cbm+10} and \citet{cbd+13}. The trend of the observations in Figure~\ref{fig_gamma} with $L_{\rm th}\gtrsim 10^{32}\,D_{\rm kpc}^2$ erg~s$^{-1}$ is roughly consistent with our expectation for a Compton-cooled bremsstrahlung model. At the lowest thermal lumnosities, however, $\Gamma$ jumps to higher values. The abrupt change is suggestive of a spectral transition to a different emission mechanism. Given the extremely low luminosities, one might consider coronal X-ray emission from the companion star \citep{br00}, but the observed spectral shapes for these observations were not consistent with coronal emission \citep{cbm+10,cbd+13}. Instead, we suggest that this may indicate a transition to synchrotron shock emission at the lowest luminosities. We discuss this further in \S\ref{sec-comps}. \subsection{The Nature of Low-$\dot M$ Accretion} \label{sec-lowMdot} As in most quiescent transient NS/LMXBs, a basic requirement for Cen X-4 is that most of the accretion flow does not reach the NS, since the inferred $\dot M_{\rm NS}$ is substantially smaller than the binary mass transfer rate $\dot M_T\sim 0.01 \dot M_E$ expected for a Roche-lobe--filing main sequence donor in a 15.1~hr binary \cite[e.g.,][]{kkb96}. There may be several mechanisms that contribute to this. First of all, the disk instability model for LMXB transients predicts that most of the accretion flow during X-ray quiescence builds up in the outer accretion disk until a thermal instability ensues, causing an outburst \citep[see][]{las01}. At low $\dot M$, the outer disk will transition into a quasi-spherical RIAF flow at $r_t$. It has previously been noted that RIAF models for quiescent NS transients require that most of the RIAF flow is somehow prevented from reaching the NS \citep{adh+98,men+99}. One way of achieving this is the ADIOS-like outflow that we discussed in \S\ref{sec-bremss2}. Another possibility is that most of the flow reaching the NS magnetosphere is centrifugally inhibited by the magnetic ``propeller effect'' \citep{is75,ukr+06}. This occurs when the magnetosphere extends beyond the corotation radius (see equation~[\ref{eq_corot}]) We define the magnetospheric radius $r_m$ as the location where the magnetic and material stresses are equal, \begin{eqnarray} r_m & = & \xi \left(\frac{\mu_m^4}{GM\dot M^2}\right)^{1/7} \\ & = & 31\ \xi\,B_8^{4/7}\,M_{1.9}^{-1/7}\,R_{10}^{10/7} \nonumber \\ & & \qquad \times\left(\frac{\dot M/\dot M_E}{0.01}\right)^{-2/7} \mbox{\rm km} , \end{eqnarray} where $\mu_m$ is the magnetic dipole moment of the NS, $\xi$ is an order unity constant that depends upon the details of the accretion flow near the magnetosphere \citep[see, e.g.,][]{pc99}, and the usual $R^{12/7}$ scaling is modified by the $R$-dependence of $\dot M_E$. In the propeller regime, $r_m>r_{\rm co}$. For ordinary thin-disk magnetic accretion, the disk extends to the magnetosphere, and the accretion is entirely shut off in this regime. However for a RIAF flow onto a millisecond pulsar, $r_m$ will generally lie inside the transition radius $r_t$, so that the flow onto the magnetosphere will be quasi-spherical. In this geometry, a small fraction of the flow is able to reach the NS despite the centrifugal barrier present in the propeller regime \citep{men+99}. Whether material is expelled in a strong or weak outflow, or else accumulates outside $r_{\rm co}$ (e.g., a ``dead'' disk), depends upon details of the disk-magnetosphere interaction \citep{st93,ds10,ds12,lru+14}. However, observationally, \citet{bcb+13} have shown that a strong propeller outflow can likely be ruled out in Cen X-4. Our observations support the conclusion of \citet{bcb+13} that low-level accretion is occurring during X-ray quiescence in Cen X-4, indicating that a small fraction of the accretion flow must eventually reach the NS. However, it remains unclear what combination of the above mechanisms ultimately controls what that fraction is. \subsection{Comparison to Other Low-$\dot M$ Systems} \label{sec-comps} After Cen X-4, the next brightest well-studied quiescent NS/LMXB transient is Aql X-1. Unlike Cen X-4, Aql X-1 has a relatively short recurrence time scale of 1--2~yr. There have been several recent studies of its quiescent emission \citep{cfh+11,ccd+14,stn+14}, all observing a soft thermal component and hard power-law component with no cutoff. \citet{stn+14} argue that, for their 2007 {\em Suzaku} observations ($L/L_E=$[3--9]$\times 10^{-5}\,M_{1.9}$ for an assumed distance of 5.2~kpc), the most appropriate model for the hard component is Comptonization in either an optically thin ($\tau_{\rm es}\simeq 0.3$), very hot ($kT_e>100$~keV) corona or an optically thick ($\tau_{\rm es}>3$), somewhat cooler ($kT_e\sim 50$~keV) corona. They do not find the same inconsistency between their measured $\tau_{\rm es}$ and radially uniform accretion that we found in equation~(\ref{eq_tau}). This is a consequence of their $\dot M$ being higher and their $\tau_{\rm es}$ being lower than in our Cen X-4 observation. However, we noted in \S\ref{sec-compt} that $\tau_{\rm es}$ and $kT_e$ vary inversely, and the high $kT_e$ values (and corresponding cutoff energies) that \citet{stn+14} fit lie above their observation bandpass. The {\em Suzaku} data are thus unable to rule out the presence of a cutoff below 50--100~keV (but still above their bandpass); this would introduce the same difficulties for a Comptonization model that we found in Cen X-4, although it would be somewhat mitigated by the higher $\dot M$. We note that our bremsstrahlung model could explain the hard component in Aql~X-1 for $kT_e\gtrsim 30$~keV. It is interesting to also compare the behavior of Cen X-4 with systems that have been observed to transition between LMXB and radio pulsar states during X-ray quiescence. The theoretical expectation is that such transitions are controlled by location of the NS magnetospheric boundary \citep{scc+94}. We can compare $r_m$ to both $r_{\rm co}$ and the light-cylinder radius, \begin{equation} r_{\rm lc} = \frac{cP}{2\pi} = 144\,P_{\rm 3ms} \mbox{\rm\ km} \,. \end{equation} For sufficiently low $\dot M$, we have $r_{\rm co}<r_m<r_{\rm lc}$ and the system will be in the propeller regime, with accretion onto the NS (mostly) shut off \citep{is75,ukr+06}. For even lower $\dot M$, we will have $r_{\rm co}<r_{\rm lc}<r_m$. In this case, the radio pulsar mechanism can turn on, with the radiation pressure of a radio pulsar wind clearing the magnetosphere and an intrabinary shock giving rise to synchrotron X-ray emission \citep{scc+94,ccm+98,bpd+01}. The $\dot M$ implied for transition to a millisecond radio pulsar state corresponds to $L_x\lesssim 10^{33}$~erg~s, where the exact value depends on details of the system and the disk-magnetosphere interaction. Indirect evidence for such transitions in NS/LMXBs during X-ray quiescence was previously reported in Aql~X-1 \citep[$L_x=6\times 10^{32}$ erg~s$^{-1}$;][]{csm+98} and SAX~J1808.4$-$3658 \citep[$L_x=5\times 10^{31}$][]{bdd+03,cdc+04,dht+08}. However, more direct evidence has been reported more recently in at least three systems. The 1.7~ms radio pulsar PSR J1023+0038 (hereafter J1023) has made two transitions between the LMXB and radio pulsar states. It is now understood to have been in an LMXB state during 2000--2001, with direct optical evidence for the presence of an accretion disk \citep{wat+09}. However, a state transition then occurred, with subsequent observations establishing the absence of an accretion disk during 2002--2013 \citep{ta05} as well as the presence of a millisecond eclipsing radio pulsar during 2007--2013 \citep{asr+09}, along with a low X-ray luminosity associated with intrabinary synchrotron shock emission \citep{akb+10,bah+11}. A second state transition was observed more recently, with the radio pulsar turning off and an accretion disk reemerging \citep{sah+13,pah+14}. In all of these observations, J1023 has remained in X-ray quiescence in the sense that a high-luminosity ($L_x\gtrsim 10^{36}$~erg~s$^{-1}$) transient X-ray outburst was not observed. However, two distinct sub-states are evident: a faint X-ray--quiescent state ($L_x\sim 10^{32}$ erg~s$^{-1}$) during which radio pulsations are seen, and a bright X-ray--quiescent state ($L_x\sim 10^{33}$ erg~s$^{-1}$) during which the radio pulsar is off. This suggests that $r_m$ has moved outside the light cylinder in the fainter state. In both sub-states, the 0.3--10~keV X-ray emission has a power-law spectrum with little or no thermal component, with $\Gamma=1.3$ in the faint state \citep{akb+10} and $\Gamma=1.69$ in the bright state \citep{pah+14}. {\em NuSTAR} observations show that these power-law spectra remain unbroken up to at least 79~keV \citep{tya+14}. During faint X-ray--quiescence, the X-ray flux shows high-amplitude modulation at the orbital period \citep{akb+10,tya+14}, similar to what is seen in the X-ray emission from some (but not all) eclipsing millisecond radio pulsars \citep[the so-called ``black widow'' and ``redback'' systems;][]{rob12,rmg+14}. During bright X-ray--quiescence, the X-ray flux shows strong, rapid flickering, with the intensity varying by an order of magnitude on time scales $<$100~s \citep{pah+14,tya+14}. A second object in which two LMXB/radio pulsar state transitions were seen is the M28 globular cluster source PSR J1824$-$2452I (hereafter M28I), a 3.9~ms radio pulsar. The radio pulsar underwent a bright transient X-ray outburst ($L_x\sim 10^{36}$ erg~s$^{-1}$) in 2013 March during which accretion-powered millisecond pulsations and a thermonuclear X-ray burst were observed, establishing the system as an NS/LMXB. The system returned to X-ray quiescence within a month, at which point radio pulsations were again detected \citep{pfb+13}. These observations demonstrate that LMXB/radio pulsar state transitions can occur on time scales as short as days. During X-ray quiescence, rapid ($\lesssim$500~s) intensity variations of nearly an order of magnitude are again seen [(0.6--4)$\times 10^{33}$ erg~s$^{-1}$], with no change in the 0.3--10~keV X-ray spectrum: an absorbed power-law with $\Gamma=1.2$ and no detectable thermal component \citep{lbh+14}. Remarkably, this rapid variability seems to toggle between two stable flux levels; \citet{lbh+14} suggest that this represents fast transitions between synchrotron shock emission and magnetospheric accretion. No orbital variability of the X-ray flux is reported in M28I. The third object in which an LMXB/radio pulsar transition was seen is XSS~J12270$-$4859 (hereafter J12270), a faint hard X-ray source associated with a relatively bright {\em Fermi} $\gamma$-ray source. During 2003--2012, the source was in a quiescent LMXB state with an absorbed power-law X-ray spectrum with $\Gamma=1.7$, no evidence for a thermal spectral component, a 0.1--10~keV luminosity of $L_x=2\times 10^{33}\,D_{\rm kpc}$~erg~s$^{-1}$, highly variable X-ray flaring, and multiwavelength evidence for the presence of an accretion disk \citep{dfb+10,dbf+13}. In late 2012, the source made a transition to a lower ($6\times 10^{31}\,D_{\rm kpc}$~erg~s$^{-1}$) luminosity state with a power-law X-ray spectrum with $\Gamma=1.2$, a thermal fraction $<$9\%, and a large-amplitude orbital modulation of the X-ray flux \citep{bph+14,bpa+14}. After this transition, 1.69~ms radio pulsations were also detected \citep{rbr14}. Cen X-4 has comparable luminosity to J1023, M28I, and J12270, and so it is presumably at least close to the regime where LMXB/radio transitions could occur. The rapid X-ray variability we observe (see \S3.1) is quite different from the large-amplitude orbital modulation seen in the low-quiescent state of J1023, but it is qualitatively similar to (although somewhat weaker than) the flickering seen in the high-quiescent states of J1023 and J12270 as well as the M28I quiescent variability. On the other hand, unlike those three sources, the X-ray spectrum of Cen X-4 has a substantial thermal fraction ($\simeq$60\%). This may indicate that more of the accretion flow reaches the surface of Cen X-4 than in the other systems. If we apply our Cen X-4 Compton-cooled bremsstrahlung model to the hard X-ray emission in J1023, M28I, and J12270 during their radio-quiet/X-ray quiescent states, then we would expect a high (50--100~keV) electron temperature and an unbroken 2--10~keV power-law X-ray spectrum, consistent with what was observed. Of course, synchrotron shock emission can also produce an unbroken power-law in the X-ray band, and we expect this mechanism to dominate in the radio pulsar state. At sufficiently low luminosity, Cen X-4 should transition into a radio pulsar state; we suggest that this may be what occurs at $L_{\rm th}\lesssim 10^{32}\,D_{\rm kpc}^2$ erg~s$^{-1}$ in Figure~\ref{fig_gamma}, and that the jump in $\Gamma$ might reflect a sharp transition from high-temperature bremsstrahlung emission to synchrotron shock emission. This is consistent with the suggestion by \citet{jgm+04} that the power-law component in quiescent NS/LMXBs arises from accretion at higher $\dot M$ and from some a different, non-accretion mechanism (e.g., synchroton shock emission) at lower $\dot M$; they used this to explain how the fractional power-law contribution to the quiescent luminosity varies with $\dot M$ in quiescent NS/LMXBs. It would be interesting to search for millisecond radio pulsations from Cen X-4 when its X-ray luminosity next drops to $\lesssim 10^{32}\,D_{\rm kpc}$~erg~s$^{-1}$.	14	3	1403.6751	The low-mass X-ray binary (LMXB) Cen X-4 is the brightest and closest (<1.2 kpc) quiescent neutron star transient. Previous 0.5-10 keV X-ray observations of Cen X-4 in quiescence identified two spectral components: soft thermal emission from the neutron star atmosphere and a hard power-law tail of unknown origin. We report here on a simultaneous observation of Cen X-4 with NuSTAR (3-79 keV) and XMM-Newton (0.3-10 keV) in 2013 January, providing the first sensitive hard X-ray spectrum of a quiescent neutron star transient. The 0.3-79 keV luminosity was 1.1× 10<SUP>33</SUP> D^2_kpc erg s<SUP>-1</SUP>, with sime60% in the thermal component. We clearly detect a cutoff of the hard spectral tail above 10 keV, the first time such a feature has been detected in this source class. We show that thermal Comptonization and synchrotron shock origins for the hard X-ray emission are ruled out on physical grounds. However, the hard X-ray spectrum is well fit by a thermal bremsstrahlung model with kT<SUB>e</SUB> = 18 keV, which can be understood as arising either in a hot layer above the neutron star atmosphere or in a radiatively inefficient accretion flow. The power-law cutoff energy may be set by the degree of Compton cooling of the bremsstrahlung electrons by thermal seed photons from the neutron star surface. Lower thermal luminosities should lead to higher (possibly undetectable) cutoff energies. We compare Cen X-4's behavior with PSR J1023+0038, IGR J18245-2452, and XSS J12270-4859, which have shown transitions between LMXB and radio pulsar modes at a similar X-ray luminosity.	false	[ "soft thermal emission", "thermal seed photons", "the first sensitive hard X-ray spectrum", "Lower thermal luminosities", "thermal Comptonization", "the hard X-ray emission", "the hard X-ray spectrum", "a similar X-ray luminosity", "keV", "the neutron star atmosphere", "a quiescent neutron star transient", "synchrotron shock origins", "the neutron star surface", "physical grounds", "unknown origin", "Cen X-4", "Compton cooling", "a thermal bremsstrahlung model", "XSS J12270", "radio pulsar" ]	5.670607	6.417522	-1
12361026	[ "Enea Romano, Antonio", "Andrés Vallejo, Sergio" ]	2015EL....10939002E	[ "Directional dependence of the local estimation of H<SUB>0</SUB> and the nonperturbative effects of primordial curvature perturbations" ]	19	[ "Yukawa Institute for Theoretical Physics, Kyoto University - Kyoto 606-8502, Japan; Department of Physics and CCTP, University of Crete - Heraklion 711 10, Greece; Department of Physics, McGill University - Montréal, QC H3A 2T8, Canada; Instituto de Fisica, Universidad de Antioquia - A.A. 1226, Medellin, Colombia", "Department of Physics and CCTP, University of Crete - Heraklion 711 10, Greece; Instituto de Fisica, Universidad de Antioquia - A.A. 1226, Medellin, Colombia" ]	[ "2012arXiv1204.0866E", "2014CQGra..31k5008R", "2015arXiv150707523B", "2016EPJC...76..216R", "2016JCAP...04..036B", "2016arXiv160406365K", "2016arXiv160904081E", "2017GReGr..49..147S", "2017IJMPD..2630011B", "2017JCAP...10..023A", "2018ApJ...854...46H", "2018GReGr..50...29K", "2018IJMPD..2750102R", "2019ApJ...881..137R", "2019JCAP...11..016C", "2019MNRAS.488.4081K", "2019arXiv191212465Y", "2020JCAP...03..023A", "2022EPJC...82..610M" ]	[ "astronomy", "physics" ]	7	[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Theory" ]	[ "1934PNAS...20..169T", "1947MNRAS.107..410B", "1986ApJ...304...15B", "1997GReGr..29..641L", "1998ApJ...493..519S", "1998ApJ...498....1W", "2005ApJ...624..463G", "2011ApJ...730..119R", "2011JCAP...10..016E", "2012EPJC...72.2242R", "2012arXiv1204.0866E", "2013ApJ...775...62K", "2013PDU.....2..166V", "2013PhRvL.110b1302B", "2013PhRvL.110x1305M", "2014A&A...571A..16P", "2014EL....10669002E", "2014EPJC...74.2780R", "2014MNRAS.438.1805W", "2014PhRvL.112v1301B" ]	[ "10.1209/0295-5075/109/39002", "10.48550/arXiv.1403.2034" ]	1403	1403.2034_arXiv.txt			14	3	1403.2034	Recent measurements of the cosmic microwave background (CMB) radiation have shown an apparent tension with the present value of the Hubble parameter inferred from local observations of supernovae, which look closer, i.e. brighter, than what is expected in a homogeneous model with a value of H<SUB>0</SUB> equal to the one estimated from CMB observations. We examine the possibility that such a discrepancy is the consequence of the presence of a local inhomogeneity seeded by primordial curvature perturbations, finding that a negative peak of the order of less than two standard deviations could allow to fit low-redshift supernovae observations without the need of using a value of the Hubble parameter different from H<SUB>0</SUB><SUP>CMB</SUP>. The type of inhomogeneity we consider does not modify the distance to the last scattering, making it compatible with the constraints of the PLANCK mission data. The effect on the luminosity distance is in fact localized around the region in space where the transition between different values of the curvature perturbations occurs, producing a local decrease, while the distance outside the inhomogeneity is not affected. Our calculation is fully relativistic and nonperturbative, and for this reason shows important effects which were missed in the previous investigations using relativistic perturbations or Newtonian approximations, because the structures seeded by primordial curvature perturbations can be today highly nonlinear, and relativist Doppler terms cannot be neglected. Because of these effects the correction to the luminosity distance necessary to explain observations is associated to a compensated structure which involves both an underdense central region and an overdense outer shell, ensuring that the distance to the last scattering surface is unaffected. Comparison with studies of local structure based on galaxy surveys analysis reveals that the density profile we find could in fact be compatible with the one obtained for the same region of sky where most of the supernovae used for the local H<SUB>0</SUB> estimation are located, suggesting a possible directional dependence which could be partially attributed to the presence of the Sloan Great Wall and hinting to the need of a more careful investigation, including a wider set of low-redshift supernovae in different regions of the sky.	false	[ "local observations", "different values", "different regions", "primordial curvature perturbations", "local structure", "CMB observations", "H<", "relativistic perturbations", "observations", "supernovae", "relativist Doppler terms", "low-redshift supernovae observations", "sky", "low-redshift supernovae", "galaxy surveys analysis", "Hubble", "one", "CMB", "Newtonian approximations", "inhomogeneity" ]	11.526824	1.578681	89
745717	[ "McCall, Marshall L." ]	2014MNRAS.440..405M	[ "A Council of Giants" ]	37	[ "Department of Physics and Astronomy, York University, Toronto, Ontario L3T 3R1, Canada" ]	[ "2014A&A...570A..13M", "2015A&A...577A.144T", "2015A&AT...29....1C", "2015ApJ...811..133D", "2015MNRAS.449.2069C", "2016IAUS..308..443C", "2017MNRAS.468.1671L", "2017MNRAS.468.1953P", "2018ApJ...864..150V", "2018RSOS....571582N", "2018arXiv180900015N", "2019MNRAS.483.2101G", "2019MNRAS.486.1964I", "2019MNRAS.487.1380G", "2020MNRAS.494.2600N", "2020MNRAS.498.2968L", "2020arXiv200101864D", "2020arXiv200900977A", "2021ApJS..256...15B", "2021MNRAS.501.3621A", "2022ApJ...935..170A", "2022ApJ...936...62L", "2022Galax..10...39B", "2022MNRAS.511.5093P", "2022arXiv221207340C", "2022icrc.confE1012B", "2023A&A...677A.169D", "2023ApJ...945..159N", "2023FrASS..1093918L", "2023JCAP...05..024A", "2023MNRAS.520L..28A", "2023MNRAS.524..631T", "2023MNRAS.526.4490P", "2023ecrs.confE..10M", "2024MNRAS.529...74A", "2024MNRAS.529.3044S", "2024arXiv240517179M" ]	[ "astronomy" ]	14	[ "galaxies: distances and redshifts", "galaxies: evolution", "galaxies: formation", "galaxies: kinematics and dynamics", "Local Group", "large-scale structure of Universe", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1953AJ.....58...30D", "1953MNRAS.113...43P", "1953ZA.....33..251K", "1955ApJ...121..161S", "1958ApJ...128..465D", "1959ApJ...130..728D", "1961BAN....16....1V", "1964ApJ...139..899D", "1970MNRAS.149..263G", "1971ApJ...170..241M", "1973A&A....24..411S", "1974ApJS...27..437T", "1974MNRAS.169..229L", "1975A&A....42..103B", "1975ApJ...202..319D", "1976PASJ...28..329O", "1977A&A....55..445F", "1977ApJ...211..324S", "1977ApJ...212..335S", "1977PASJ...29..567O", "1978ApJ...220...98S", "1979A&A....74...73S", "1979AJ.....84..284D", "1979AN....300..181K", "1979ApJ...229...91T", "1979ApJ...229..509R", "1980A&A....90..123V", "1980ApJ...239...54P", "1980ApJ...239..783D", "1980MNRAS.190..689N", "1980MNRAS.191..169N", "1981ApJ...247..473P", "1982A&A...105...76M", "1982ApJS...49..515D", "1983A&A...118...33K", "1983ApJ...265..664D", "1984A&A...132...20S", "1984A&A...141..309B", "1985ApJ...289...81R", "1985ApJ...289..129S", "1985ApJS...58..107C", "1985Natur.317...44C", "1986A&A...167...34B", "1986A&AS...63..323H", "1986A&AS...66..505W", "1986AJ.....91..777B", "1986MNRAS.218..297W", "1987A&AS...69..311W", "1987AJ.....94.1519C", "1988A&A...196...26K", "1988ARA&A..26..509B", "1988ApJS...66..233B", "1989ApJ...339...53C", "1989ApJ...344..704J", "1989ApJ...344..715C", "1990A&A...236...33H", "1990AJ.....99.1813C", "1990AJ....100..641C", "1990AJ....100.1468P", "1991A&A...244....8T", "1991AJ....101..447P", "1991AJ....101.1231C", "1991ApJ...383..487C", "1992A&A...253..335K", "1992AJ....103..793S", "1992AN....313..189S", "1992AN....313..329S", "1992ApJ...387...47P", "1992MNRAS.258..334S", "1993A&A...269...15V", "1993A&A...270...29D", "1993A&A...274...69C", "1993ApJ...414..463H", "1993PASJ...45..139S", "1994A&A...285..833R", "1994A&AS..106..451D", "1994AJ....108.1610M", "1994ApJ...430..142R", "1995AJ....109.1592R", "1995AJ....109.2038I", "1995ApJ...438..155R", "1995ApJ...439..163A", "1995ApJ...448L..13B", "1995ApJ...449..592H", "1995PhR...256..157M", "1996A&A...306....9S", "1996A&A...309..687B", "1996A&A...314...32B", "1996A&AS..118..111H", "1996AJ....111..735A", "1996AJ....112..457O", "1996AJ....112.2471T", "1996ApJ...458..455F", "1996ApJ...465...79S", "1996ApJ...466..135H", "1997ApJ...479..231F", "1997ApJ...491..140S", "1997MNRAS.286..771E", "1997PASJ...49...17S", "1997PASJ...49..279K", "1998ASPC..134..483K", "1998ASPC..147..259M", "1998ApJ...495L..47H", "1998ApJ...500..525S", "1999A&A...342...87O", "1999A&AS..136..509H", "1999A&AS..139..491M", "1999AJ....117..839S", "1999AJ....117..855H", "1999AJ....117.2666N", "1999AJ....118.1209F", "1999AJ....118.2184H", "1999ApJ...526..599S", "1999ApJS..124...33B", "1999MNRAS.302..649J", "1999MNRAS.303..495E", "1999PASP..111...63F", "2000A&A...356..827V", "2000AJ....119.1157G", "2000AJ....120..763P", "2000ASPC..218....1M", "2000ApJ...529..698S", "2000ApJS..127...39C", "2000ApJS..128..431F", "2000ApJS..128..461M", "2001A&A...370..765V", "2001A&A...372..463O", "2001A&A...374..394V", "2001AJ....121.2557D", "2001ApJ...546..681T", "2001ApJ...553...47F", "2001ApJ...554..104P", "2001ApJ...563..694V", "2001ApJS..132..129M", "2001BaltA..10..413B", "2001MNRAS.327.1004B", "2002A&A...383..125K", "2002A&A...384..371S", "2002A&A...385...21K", "2002A&A...390..863S", "2002A&A...392..473S", "2002A&A...394..769G", "2002AJ....123..244K", "2002AJ....123.1892C", "2002AJ....123.3124F", "2002AJ....124..839S", "2002ApJ...569L..69M", "2002ApJ...577...31C", "2002MNRAS.329..513D", "2003A&A...398..467K", "2003A&A...398..479K", "2003A&A...399...51G", "2003A&A...400..451G", "2003A&A...404...93K", "2003A&A...412...45P", "2003AJ....125..525J", "2003ApJ...587..672F", "2003ApJ...590..256T", "2003ApJ...594..279G", "2003ApJ...596...19K", "2003ApJS..146..299E", "2003IAUS..209..583M", "2003MNRAS.338..465A", "2004A&A...415...63M", "2004A&A...415..889C", "2004A&A...421..433M", "2004A&A...423..925G", "2004AJ....127.1472B", "2004AJ....127.2031K", "2004AJ....128...16K", "2004AJ....128.2144M", "2004ApJ...605..183W", "2004ApJ...608...42S", "2004ApJ...613..914H", "2004ApJ...614..167C", "2004MNRAS.347..691P", "2004MNRAS.347..968P", "2004MNRAS.350..243M", "2004MNRAS.352..721E", "2005A&A...432..475D", "2005AJ....129..178K", "2005AJ....129.1331S", "2005AJ....130..406S", "2005AJ....130.1593V", "2005ApJ...618..569M", "2005ApJ...627..647B", "2005ApJ...628L..33M", "2005ApJ...631..262R", "2005ApJ...633..810M", "2005AstL...31..299K", "2005MNRAS.356..979M", "2005MNRAS.360.1201H", "2005MNRAS.361..330R", "2006AJ....131.1163S", "2006AJ....131.1361K", "2006ApJ...641L.109C", "2006ApJ...652.1133M", "2006ApJS..162...49H", "2006MNRAS.367..449H", "2006MNRAS.367..469D", "2006MNRAS.367..815N", "2006MNRAS.372.1149F", "2006MNRAS.373.1265O", "2007A&A...463..427P", "2007AJ....133..504K", "2007AJ....134..494W", "2007AJ....134.1019O", "2007ApJ...654..186F", "2007ApJ...655..814F", "2007ApJ...661..815R", "2007ApJS..169..225E", "2007MNRAS.379..418C", "2008AJ....135.1900A", "2008AJ....136..773F", "2008AJ....136.1866K", "2008AJ....136.2648D", "2008ARA&A..46..201S", "2008ApJ...672..266G", "2008ApJ...676..184T", "2008ApJ...683..630H", "2008ApJ...684.1143X", "2008ApJ...686L..75M", "2008ApJ...686L..79G", "2008MNRAS.384..943N", "2008MNRAS.389...63C", "2008MNRAS.389.1001L", "2009ApJ...694.1331M", "2009ApJ...700..137R", "2009PASJ...61..227S", "2010ApJ...725.2087P", "2010MNRAS.403.1829S", "2010Natur.465..565P", "2011ApJ...730..119R", "2011MNRAS.414.2446M", "2011MNRAS.416..509H", "2012ApJ...752...76R", "2012ApJ...753....8V", "2012MNRAS.421..847C", "2012MNRAS.421L.137L", "2012MNRAS.422.1732D", "2013ApJ...766..120C", "2013ApJ...775...13H", "2013Natur.493...62I", "2014A&A...571A..16P", "2014MNRAS.437.2111V" ]	[ "10.1093/mnras/stu199", "10.48550/arXiv.1403.3667" ]	1403	1403.3667.txt	%Done Galaxies are organized into an expanding cosmic web of filamentary and sheet-like structures bounding volumes which are largely devoid of matter. However, very little is known observationally about the structure of structures and its linkages to galaxy evolution because it is difficult to constrain accurate relative positions of constituents from a distant vantage point. The Local Sheet, a structure of which we are a part, offers an opportunity for advancement owing to our perspective from within and the proximity to measure reliable distances to members directly. Any study of local structure must start with a volume-limited sample of galaxies. Efforts to construct such a sample began with the definition of the Local Volume \citep{kra79a,huc86a,sch92a}, %FINAL which in the rendition initiating this work (the Local Volume Catalog, or LVC -- \citealt{kar04a}) contains all known galaxies either with distances less than $10 \, \rm Mpc$ %FINAL or with radial velocities less than $550 \, \rm km \, s^{-1}$ %FINAL with respect to the Local Group (a Hubble flow distance of $7.7 \, \rm Mpc$). %FINAL Within the Local Volume, the Milky Way, Andromeda, and the smaller companions which comprise the Local Group reside in a layer of galaxies, mostly dwarfs, which has an apparent thickness of about $1.5 \, \rm Mpc$ \citep{sch92b,pee93a,pee01a,kar04a,kar05a,tul08a,fin12a}. %FINAL At various times, the layer has been referred to as the ``Local Cloud'' \citep{vau75a}, the ``Coma-Sculptor Cloud'' \citep{tul87a}, %FINAL the ``local plane'' \citep{pee93a}, %FINAL the ``local filament'' \citep{kly03a}, %FINAL the ``Local pancake'' \citep{kar04a}, %FINAL and the ``Local Sheet'' \citep{pee01a,tul08a,pee10a}. %FINAL At a certain level, it is the proximate manifestation of the Local Supercluster, whose existence was in fact established in part using the most luminous members of the layer \citep{vau53a}. %FINAL However, models of the local velocity field seem to require that the Local Group be housed in a flattened body of galaxies distinct from the Local Supercluster \citep{kly03a}. %FINAL Indeed, it has been argued that the supergalactic arrangement of nearby groups in the plane of the sky is evidence for such a body \citep{vau75a}. %FINAL Also, peculiar velocities of galaxies show a sharp discontinuity at a distance of about $7 \, \rm Mpc$ \citep{tul08a}. %FINAL Studies of local structure are traditionally anchored to the supergalactic coordinate system, but whether or not this is the appropriate framework to adopt has not been examined thoroughly. Any local flattened structure distinct from the Local Supercluster ought to be traced most reliably by its most {\color{black} luminous} members, because they pinpoint the location of the largest concentrations {\color{black} of dark matter}. Consequently, to isolate such a structure {\color{black} and elucidate its character}, this paper focuses on carefully mapping the distribution and properties of luminous spiral and elliptical galaxies in the Local Volume. {\color{black} The framework is vital for guiding studies of the dwarf population locally} \citep{fin12a}, results from which will be presented separately.	%Done {\color{black} Properties of the Local Sheet and its Council of Giants are summarized in Table~\ref{tbl_summary}. } This study suggests that a structure with the geometry of the Sheet was instrumental in guiding the formation and evolution of constituent galaxies. It also suggests that a binary, or the precursor of it, can influence the angular momentum acquired by neighbouring galaxies. It is unlikely that a randomly dispersed collection of galaxies could have agglomerated into a structure as cold as the Sheet in a way which could generate an interacting pair of galaxies {\color{black} near the middle of a ring of galaxies with opposing ellipticals and ordered spins}. Indeed, modern cosmological simulations reveal galaxies developing from dark cores fed by flows of gas along pre-existing filaments of dark matter \citep{dan12a}. Unfortunately, there is only one Local Sheet. Further insights into the formation and evolution of the Local Sheet, and particularly guidance on the interplay between the Local Group and the Local Sheet, will require the identification of like systems in the greater Universe.	14	3	1403.3667	Distances and near-infrared luminosities of the brightest galaxies in the Local Volume have been re-evaluated in order to gain a fully homogeneous collection of data for elucidating the framework of the Local Sheet and its relevance to Local Group evolution. It is demonstrated that the Local Sheet is both geometrically and dynamically distinct from the Local Supercluster and that the evolution of the Sheet and Local Group were probably interconnected. The Sheet is inclined by 8° with respect to the Local Supercluster, and the dispersion of giant members about the mid-plane is only 230 kpc. A `Council of Giants' with a radius of 3.75 Mpc encompasses the Local Group, demarcating a clear upper limit to the realm of influence of the Local Group. The only two giant elliptical galaxies in the Sheet sit on opposite sides of the Council, raising the possibility that they have somehow shepherded the evolution of the Local Group. The position vector of the Andromeda galaxy with respect to the Milky Way deviates only 11° from the Sheet plane and only 11° from the projected axis of the ellipticals. The Local Group appears to be moving away from a ridge in the potential surface of the Council on a path parallel to the elliptical axis. Spin directions of the giants in the Council are distributed over the sky in a pattern which is very different from that of giants beyond, possibly in reaction to the central mass asymmetry that developed into the Local Group. By matching matter densities of Group and Council giants, the edge of the volume of space most likely to have contributed to the development of the Local Group is shown to be very close to where gravitational forces from the Local Group and the Council balance. The boundary specification reveals that the Local Sheet formed out of a density perturbation of very low amplitude (∼10 per cent), but that normal matter was incorporated into galaxies with relatively high efficiency (∼40 per cent). It appears that the development of the giants of the Local Sheet was guided by a pre-existing flattened framework of matter.	false	[ "Local Group", "Local Group evolution", "Local Supercluster", "The Local Group", "the Local Group", "Sheet", "Council", "the Local Sheet", "Group and Council giants", "the Sheet and Local Group", "cent", "Group and Council", "matter", "the Local Supercluster", "the Local Volume", "galaxies", "giants", "plane", "the Sheet plane", "opposite sides" ]	10.732628	6.20607	-1
746013	[ "Shivvers, Isaac", "Bloom, Joshua S.", "Richards, Joseph W." ]	2014MNRAS.441..343S	[ "The highly eccentric detached eclipsing binaries in ACVS and MACC" ]	9	[ "Astronomy Department, University of California, Berkeley, CA 94720, USA", "Astronomy Department, University of California, Berkeley, CA 94720, USA", "Astronomy Department, University of California, Berkeley, CA 94720, USA" ]	[ "2015ApJS..217...28Y", "2016PASP..128g4201K", "2018ApJS..235...41K", "2021A&A...652A..81Z", "2022MNRAS.510.2448H", "2022NewA...9701875K", "2022RMxAA..58....3C", "2023A&A...670A..39Z", "2024arXiv240512136M" ]	[ "astronomy" ]	10	[ "techniques: spectroscopic", "catalogues", "binaries: eclipsing", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1879Obs.....3...79D", "1959ASPL....8...81K", "1971ApJ...166..605W", "1974A&A....31..129S", "1975A&A....41..329Z", "1976Ap&SS..39..447L", "1977A&A....57..383Z", "1978ApJ...224..953S", "1979AJ.....84.1511T", "1979ApJ...234.1054W", "1980ARA&A..18..115P", "1981A&A....99..126H", "1982ApJ...263..835S", "1985Ap&SS.114..259G", "1987PASP...99.1214V", "1988ApJ...324L..71T", "1989A&A...223..112Z", "1991A&ARv...3...91A", "1994ApJ...420..806Z", "1995A&A...299..724C", "1996FCPh...16..337T", "1996PASP..108..500B", "1997A&A...318..187C", "2000A&AS..141..371G", "2002A&A...385.1095P", "2002A&A...391..369K", "2002ApJS..141..503N", "2003A&A...405..677N", "2004AcA....54..153P", "2004MNRAS.351.1277S", "2005ApJ...620..970M", "2005ApJ...628..411D", "2005ApJ...628..426P", "2005ApJS..159..141V", "2008ApJS..175..191P", "2010A&ARv..18...67T", "2011AJ....142...60V", "2011ApJ...733...10R", "2012ApJS..203...32R", "2013A&A...556A..56D", "2013A&A...556A.138F", "2013ApJ...763L...2D", "2013MNRAS.434..925H", "2013PASP..125.1336T" ]	[ "10.1093/mnras/stu578", "10.48550/arXiv.1403.5564" ]	1403	1403.5564_arXiv.txt	Detached eclipsing binaries (DEBs) have long served as powerful astronomical tools, allowing astronomers to derive the fundamental parameters of distant stars \citep[e.g.][]{kopal59,popper80,gimenez85,anderson91,Torres}. The most important results of these binary star studies are facilitated by the assumption that stars in non-interacting (`detached') orbits evolve as if they were single stars, an assumption which allows us to map our empirical understanding of detached binary systems directly onto single stars everywhere. This assumption generally appears to be robust, but there are several outstanding problems in our understanding of the evolutionary differences between stars in binary pairs and those that have evolved singly. One of the most important gaps in our understanding has to do with the history of orbital circularization in binary stars and just how important each star's individual tidal interaction history is. Many binary stars display orbital eccentricities very near zero. The primordial distribution of eccentricities formed from binary star formation is only poorly understood, but it is expected to be much wider than that exhibited in evolved stars. The low eccentricities of evolved binaries is likely a result of tidal interactions between the two stars conspiring to circularize the orbit \citep[e.g.][and citations within]{mazeh}. To first order, any primordial eccentricity is expected to decay exponentially on a characteristic timescale ($t_{circ}$) which depends strongly upon internal structure and orbital separation \citep[tidal circularization is a complex topic that has featured in astrophysical discussions for many years; e.g.][]{darwin}. Several different processes likely contribute to this decay, including dynamical and equilibrium tide effects \citep[e.g.][]{zahn75, zahn77, hut} and hydrodynamical effects \citep[e.g.][]{tassoul88, tassoul96}, depending on the masses, sizes, structure and separation of the binary components. Final estimates of the circularization timescale can vary widely \citep[e.g.][]{zahn89}, and analyses of systems in the literature indicate problems with our understanding \citep[e.g.][]{claret1, claret2, north, meibom}. However, it is abundantly clear that most binary stars undergo some variation of circularization for some amount of their evolutionary history. This raises an important question: how does this process affect the binary components themselves, and could it introduce a systematic difference between the properties of stars in detached binaries and singly-evolved stars, thereby affecting our measurements of the Mass-Radii or Mass-Luminosity relations for all stars? To answer this question empirically we set out to identify and characterize binary stars in highly-eccentric orbits. Whether because they are young or because their $t_{circ}$ is very large, these systems have (on average) undergone less of the circularization process than their $e = 0.0$ counterparts, and any systematic change wrought by that process may be measurable by comparing the properties of binary stars in high-$e$ orbits to those in circular orbits. With the advent of very large photometric data sets and new techniques for classification the number of known DEBs has skyrocketed \citep[e.g.][]{anderson91,debil,Torres,MACC}, but the number of those that have well-measured masses and radii is quite small. A recent review by \citet{Torres} presented the currently-known DEB systems with physical system parameters determined to an accuracy of $\pm 3 \%$ or better, of which only $15$ have $e \gtrsim 0.1$. Radial velocity (RV) curves are required but they are time-intensive and expensive to produce and most of the high-$e$ systems found by modern surveys are relatively faint, making it difficult to obtain the high-resolution spectra necessary for accurate RVs. However, the All-Sky Automated Survey Catalog of Variable Stars (ASAS/ACVS) photometric database probes magnitude ranges amenable to RV followup, has a long baseline of observations with a relatively high cadence (producing the well-sampled light curves needed to identify and characterize eclipsing binaries) and was recently re-analyzed with all sources re-classified through modern machine-learning techniques by \citet{MACC}, yielding the Machine-learned ASAS Classification Catalog (MACC). DEBs are identified within MACC, though their degree of orbital eccentricity is not. Helpfully, the techniques used by MACC enable targeted searches for anomalous sources and for rare classes of objects. We performed a search for highly-eccentric DEBs ($e \gtrsim 0.1$) in the MACC, yielding 106 bright (V$<$12\,mag) systems, most of which are identified here for the first time. We also present new high-resolution RV curves with modeled physical parameters for six of these objects. In \S \ref{sec:data} we describe our photometric and spectroscopic data; in \S \ref{sec:finding} we describe our search for eccentric eclipsing binaries and present the highly-eccentric DEBs in MACC; in \S \ref{sec:RV} we present the new RV curves and physical parameters for six systems; in \S \ref{sec:conclusion} we discuss our results.	\label{sec:conclusion} Figure \ref{fig:eccents} displays the best-fit eccentricities for these sources alongside the results of previous studies of other high-$e$ binary systems. None of the systems presented here appear to be extreme outliers, though several do exhibit a remarkably high $e$ for their period. Figure \ref{fig:RvM} plots the radii and masses of the 12 stars presented in Table~\ref{tab:params} alongside the results of previous studies. Note that the population of stars included on that plot is heavily affected by several observational biases and any population studies should be undertaken with caution. For example, only systems with relatively strong secondary eclipses are identified as DEBs by surveys like this one. For this situation to arise a system with two main-sequence stars must have similar effective temperatures and therefore similar masses, which likely explains the nearly twin properties exhibited by some of the systems presented in this paper. In contrast, binary systems with unequal masses will preferentially be detected only after the more massive primary star has moved away from the main sequence and expanded enough to reach an effective temperature comparable to that of the secondary star; this may help to explain the relatively large radii exhibited by a few of these stars presented here. A complete analysis is beyond the scope of this paper; we encourage further studies to constrain the effect of tidal cicularization on stellar binary components, and look forward to the important results coming from modern photometric surveys like Kepler \citep[e.g.][]{deboss13,frandsen13,hambleton} and the Palomar Transient Factory \citep[e.g.][]{Eyken11,Hamren11,Prince13}, as well as the next generation of synoptic photometric surveys. We have outlined the process and results of a targeted search for highly-eccentric DEBs using the MACC probabilistic catalog and we presented 106 such systems with eccentricity estimates, most of which are listed here for the first time. We presented new RV curves and modeled masses and radii for 6 of these systems, each of them double-lined spectroscopic binaries. We publish these data to facilitate further studies of orbital evolution through tidal dissipation. We also present this project as an example of how modern large-scale datasets can be immediately used to address outstanding problems in a wide range of astronomical subfields. With limited funds and resources for new data acquisition, taking full advantage of the data already in hand and identifying the sources most worthy of further study is vitally important. As we assemble these large-scale photometric datasets, we must also become more adept at understanding the data and locating rare objects of interest buried in them. \begin{figure} \centering \includegraphics[width=\linewidth]{RvM.eps} \caption{Radii and masses for the 12 stars presented in Table~\ref{tab:everything} and \citet{Torres}. Grey (circle): $e<0.1$ systems from \citet{Torres}; blue (diamond): $e>0.1$ systems from \citet{Torres}; red (square): new systems presented in this work. Note that errorbars are often smaller than the plotted symbol. The dashed line shows a theoretical zero-age main sequence for stars of solar metallicity \citep{girardi}. The 12 stars added here to this plot appear to fall within the normal ranges, though a complete analysis is beyond the scope of this paper.\label{fig:RvM}} \end{figure} \FloatBarrier	14	3	1403.5564	Next-generation synoptic photometric surveys will yield unprecedented (for the astronomical community) volumes of data and the processes of discovery and rare-object identification are, by necessity, becoming more autonomous. Such autonomous searches can be used to find objects of interest applicable to a wide range of outstanding problems in astronomy, and in this paper we present the methods and results of a largely autonomous search for highly eccentric detached eclipsing binary systems in the Machine-learned All-Sky Automated Survey Classification Catalog. 106 detached eclipsing binaries with eccentricities of e ≳ 0.1 are presented, most of which are identified here for the first time. We also present new radial-velocity curves and absolute parameters for six of those systems with the long-term goal of increasing the number of highly eccentric systems with orbital solutions, thereby facilitating further studies of the tidal circularization process in binary stars.	false	[ "Such autonomous searches", "binary stars", "highly eccentric detached eclipsing binary systems", "outstanding problems", "objects", "further studies", "orbital solutions", "absolute parameters", "astronomy", "necessity", "highly eccentric systems", "the tidal circularization process", "interest", "the Machine-learned All-Sky Automated Survey Classification Catalog", "results", "first", "data", "volumes", "discovery and rare-object identification", "106 detached eclipsing binaries" ]	6.958603	11.169543	81
483177	[ "Del Moro, A.", "Mullaney, J. R.", "Alexander, D. M.", "Comastri, A.", "Bauer, F. E.", "Treister, E.", "Stern, D.", "Civano, F.", "Ranalli, P.", "Vignali, C.", "Aird, J. A.", "Ballantyne, D. R.", "Baloković, M.", "Boggs, S. E.", "Brandt, W. N.", "Christensen, F. E.", "Craig, W. W.", "Gandhi, P.", "Gilli, R.", "Hailey, C. J.", "Harrison, F. A.", "Hickox, R. C.", "LaMassa, S. M.", "Lansbury, G. B.", "Luo, B.", "Puccetti, S.", "Urry, M.", "Zhang, W. W." ]	2014ApJ...786...16D	[ "NuSTAR J033202-2746.8: Direct Constraints on the Compton Reflection in a Heavily Obscured Quasar at z ≈ 2" ]	28	[ "Department of Physics, Durham University, South Road, Durham, DH1 3LE, UK", "Department of Physics, Durham University, South Road, Durham, DH1 3LE, UK; Department of Physics & Astronomy, University of Sheffield, Sheffield, S3 7RH, UK", "Department of Physics, Durham University, South Road, Durham, DH1 3LE, UK", "INAF-Osservatorio Astronomico di Bologna, Via Ranzani 1, I-40127 Bologna, Italy", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, 306, Santiago 22, Chile; Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO 80301, USA", "Departamento de Astronomía, Universidad de Concepción, Casilla 160-C, Concepción, Chile", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Mail Stop 169-221, Pasadena, CA 91109, USA", "Department of Physics and Astronomy, Dartmouth College, 6127 Wilder Laboratory, Hanover, NH 03755, USA; Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "National Observatory of Athens, Institute of Astronomy, Astrophysics, Space Applications and Remote Sensing, Metaxa & Pavlou St., 15236 Penteli, Greece", "INAF-Osservatorio Astronomico di Bologna, Via Ranzani 1, I-40127 Bologna, Italy; Dipartimento di Fisica e Astronomia, Università di Bologna, via Berti Pichat 6/2, I-40127 Bologna, Italy", "Department of Physics, Durham University, South Road, Durham, DH1 3LE, UK", "Center for Relativistic Astrophysics, School of Physics, Georgia Institute of Technology, Atlanta, GA 30332, USA", "Cahill Center for Astrophysics, 1216 East California Boulevard, California Institute of Technology, Pasadena, CA 91125, USA", "Space Sciences Laboratory, University of California, Berkeley, CA 94720, USA", "Department of Astronomy and Astrophysics, 525 Davey Lab, The Pennsylvania State University, University Park, PA 16802, USA; Institute for Gravitation and the Cosmos, The Pennsylvania State University, University Park, PA 16802, USA", "DTU Space-National Space Institute, Technical University of Denmark, Elektrovej 327, DK-2800 Lyngby, Denmark", "Lawrence Livermore National Laboratory, Livermore, CA 94550, USA", "Department of Physics, Durham University, South Road, Durham, DH1 3LE, UK", "INAF-Osservatorio Astronomico di Bologna, Via Ranzani 1, I-40127 Bologna, Italy", "Columbia Astrophysics Laboratory, 550 W 120th Street, Columbia University, NY 10027, USA", "Cahill Center for Astrophysics, 1216 East California Boulevard, California Institute of Technology, Pasadena, CA 91125, USA", "Department of Physics and Astronomy, Dartmouth College, 6127 Wilder Laboratory, Hanover, NH 03755, USA", "Yale Center for Astronomy & Astrophysics, Yale University, Physics Department, P.O. Box 208120, New Haven, CT 06520-8120, USA", "Department of Physics, Durham University, South Road, Durham, DH1 3LE, UK", "Department of Astronomy and Astrophysics, 525 Davey Lab, The Pennsylvania State University, University Park, PA 16802, USA; Institute for Gravitation and the Cosmos, The Pennsylvania State University, University Park, PA 16802, USA", "ASI-Science Data Center, via Galileo Galilei, I-00044 Frascati, Italy; INAF-Osservatorio Astronomico di Roma, via Frascati 33, I-00040 Monteporzio Catone, Italy", "Yale Center for Astronomy & Astrophysics, Yale University, Physics Department, P.O. Box 208120, New Haven, CT 06520-8120, USA", "NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA" ]	[ "2014ApJ...792..108B", "2014ApJ...792..117G", "2014ApJ...794..102S", "2014ApJ...794..111B", "2014MNRAS.443.1999B", "2015A&ARv..23....1B", "2015ApJ...807..129S", "2015ApJ...807..149K", "2015ApJ...808..184M", "2015ApJ...808..185C", "2015ApJ...809..115L", "2015ApJ...812..116B", "2015ApJ...815...66A", "2015MNRAS.451.1892A", "2016ApJ...831...76F", "2016ApJ...831..185H", "2016MNRAS.456.2105D", "2017A&ARv..25....2P", "2017ApJ...835..105R", "2017ApJ...836...99L", "2017ApJ...846...20L", "2017ApJ...849...53H", "2017ApJ...849...57D", "2017ApJS..228....2L", "2017ApJS..232....8L", "2017NewAR..79...59X", "2018ApJ...854...33Z", "2021PASA...38...33B" ]	[ "astronomy" ]	3	[ "galaxies: active", "infrared: galaxies", "quasars: general", "quasars: individual: NuSTAR J033202-2746.8", "X-rays: galaxies", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1955ApJ...121..161S", "1962PhRvL...9..439G", "1978AnSta...6..461S", "1979ApJ...228..939C", "1990ARA&A..28..215D", "1991MNRAS.249..352G", "1992xrb..work...44G", "1993ARA&A..31..473A", "1994ApJ...427..134W", "1994MNRAS.267..743G", "1994MNRAS.268..405N", "1995A&A...296....1C", "1995ApJ...448L..81W", "1995MNRAS.273..837M", "1995PASJ...47L...5G", "1995PASP..107..375O", "1996MNRAS.280..823M", "1998A&A...333..411V", "1998ARA&A..36..189K", "1998ApJ...493..650M", "1999A&A...349L..73V", "1999ApJ...520..124G", "1999ApJ...526...60S", "1999MNRAS.309..969H", "2000ApJ...537..654E", "2000MNRAS.316..234R", "2001A&A...377..442W", "2001ApJ...556L..91S", "2002A&A...393..425M", "2002ApJ...573L..81L", "2002ApJS..139..369G", "2002PASJ...54..327K", "2003AJ....126..539A", "2003ApJ...598..886U", "2003ApJS..148..175S", "2004A&A...414..767K", "2004A&A...416..901C", "2004A&A...418..465L", "2004A&A...421..913W", "2004ApJS..155...73Z", "2004ApJS..155..271S", "2005A&A...444...79M", "2005ApJ...630..115T", "2005ApJS..161...21L", "2005MNRAS.357.1281W", "2005MNRAS.358..693C", "2005MNRAS.364..195P", "2006A&A...448..499B", "2006A&A...451..457T", "2006A&A...455..173P", "2006ApJ...639..740B", "2006ApJ...642...96E", "2006ApJ...645...95H", "2006MNRAS.370.1893H", "2006SPIE.6270E..1VF", "2007A&A...463...79G", "2007A&A...467..519P", "2007A&A...475..775K", "2007ApJ...660L..43N", "2007ApJ...663..204B", "2007ApJ...664L..79U", "2007ApJS..172....1S", "2007ApJS..172...29H", "2007MNRAS.382..194N", "2007MNRAS.382.1005G", "2008ApJ...687..835A", "2008ApJ...689..666A", "2008ApJS..179...19L", "2008ApJS..179...71K", "2008ApJS..179...95M", "2009A&A...493..445D", "2009A&A...502..457G", "2009ApJ...693..447F", "2009ApJ...696..110T", "2009ApJ...698..740T", "2009MNRAS.394L.121M", "2010ApJ...714.1582B", "2010ApJ...717..787C", "2010ApJS..187..560L", "2010ApJS..189..270C", "2011A&A...526L...9C", "2011A&A...528A..35M", "2011A&A...530A..42C", "2011A&A...532A..90L", "2011A&A...533A.119E", "2011A&A...536A.105V", "2011ApJ...726...20M", "2011ApJ...728...58B", "2011ApJ...736...56B", "2011ApJS..195...10X", "2011MNRAS.414.1082M", "2012A&A...546A..84I", "2012A&A...546A..98A", "2012ApJ...758...90S", "2012ApJ...758..129X", "2012ApJS..201...34B", "2012ApJS..203...15B", "2012MNRAS.419...95M", "2012MNRAS.423.3360Y", "2012MNRAS.424.1614O", "2013A&A...549A..59D", "2013A&A...553A.132M", "2013A&A...555A..42R", "2013A&A...555A..43G", "2013A&A...556A.114C", "2013ApJ...770..103H", "2013ApJ...773..125A", "2013ApJS..205...13M", "2013MNRAS.429.2407H", "2013Natur.494..449R", "2014ApJ...792...48W", "2014MNRAS.437.2845B" ]	[ "10.1088/0004-637X/786/1/16", "10.48550/arXiv.1403.2491" ]	1403	1403.2491_arXiv.txt	Many studies in the past 50 years have been devoted to understanding the origin of the X-ray background (XRB) radiation since its discovery in the early 1960's \citep{giacconi1962}. It is now clear that this radiation is due to the emission from individual X-ray sources, with Active Galactic Nuclei (AGN) being the main contributors to the overall XRB emission. The XRB spectrum, as measured from several past and current X-ray missions (e.g., {\it HEAO-1 A2}, \citealt{gruber1992,gruber1999}; {\it BeppoSAX}, \citealt{vecchi1999}; {\it ASCA}, \citealt{gendreau1995,kushino2002}; {\it Swift}-BAT, \citealt{ajello2008}), peaks at $\approx20-30$ keV. Many studies infer that a large population of heavily obscured and Compton-thick AGN (with column densities of $N_{\rm H}>10^{24}$ cm$^{-2}$) are needed to produce this peak (e.g., \citealt{comastri1995,worsley2005,treister2005,ballantyne2006,gilli2007,treister2009}). However, there are still uncertainties regarding the relative contribution of obscured and Compton-thick AGN populations to the XRB spectrum. Deep X-ray surveys from the \ch\ and \xmmn\ observatories (e.g., \citealt{alexander2003, hasinger2007, comastri2011, xue2011, ranalli2013}) have provided the best constraints on the source population dominating the X-ray emission at $E<10$ keV, allowing us to resolve $\approx70-90$\% of the XRB at energies $E\approx0.5-10$ keV (e.g., \citealt{worsley2005, hickox2006, xue2012}). However, due to their $\approx0.1-10$ keV bandpass, \ch\ and \xmmn\ cannot provide a clear indication of the population dominating at higher energies ($E>10$ keV); this also means that these observatories are relatively insensitive to the identification of the most heavily obscured AGN, where the low-energy emission is suppressed by large column density gas. On the other hand, until recently the X-ray telescopes sensitive at energies $E>10$ keV (e.g., {\it Swift}-BAT, {\it INTEGRAL} and {\it Suzaku}) have yielded direct constraints of only $\approx1-2$\% of the hard X-ray population contributing to the XRB emission at its peak (e.g., \citealt{krivonos2007,ajello2008,bottacini2012}) due to their inherently high background levels. As a consequence, there are still large uncertainties on the predictions from the models of the XRB, which vary significantly depending on the assumed distribution of absorbing column densities, the intrinsic X-ray spectral properties, and the fraction of heavily obscured and Compton-thick AGN (e.g., \citealt{gilli2007,treister2009,ballantyne2011,akylas2012}). Today, great improvements can be made in measuring the composition of the XRB thanks to the {\it Nuclear Spectroscopic Telescope Array} (\nus; \citealt{harrisonf2013}). \nus\ is the first high-energy orbiting observatory ($E\approx3-79$ keV) equipped with focussing optics, which make this satellite $\approx$2 orders of magnitude more sensitive than the previous-generation hard X-ray ($E>10$ keV) observatories, with one order of magnitude higher angular resolution. With its unique characteristics \nus\ allows us to: 1) identify sources almost independently from their level of obscuration (at least for column densities $N_{\rm H}\lesssim10^{25}$ cm$^{-2}$), therefore overcoming the limitations of the lower-energy observatories currently available; 2) measure the composition of the XRB at its peak energies, providing direct constraints on the contribution from different AGN populations; 3) characterize the broad-band X-ray spectra of AGN, removing ambiguities on the source properties (which are often present when only $E<10$ keV spectra are available), yielding unprecedented constraints on the spectral models (e.g., \citealt{risaliti2013}) even for heavily obscured and Compton-thick AGN out to high redshift ($z\approx2$). % In this paper we investigate the case of NuSTAR J033202--2746.8, which is detected in the \nus\ observations of Extended \ch\ Deep Field-South (E-CDF-S; Mullaney et al., in prep.). NuSTAR J033202--2746.8 is only significantly detected at $E>8$ keV, showing the hardest band ratio among all of the \nus\ sources detected in the E-CDF-S field so far. This suggests that the source is heavily obscured. Although previously studied at low energies with \ch\ and \xmmn, the X-ray spectrum of this source has never been accurately characterized. Indeed, with the \nus\ data that allow us to constrain its spectral parameters over a broad energy range, we identify a significant reflection component contributing to the spectrum at high energies. Furthermore, using deep infrared (IR) data available in this field, we independently infer the intrinsic power of the AGN, as well as characterize its host galaxy by means of detailed Spectra Energy Distribution (SED) analysis. % The paper is organized as follows: in section \ref{data} we briefly describe the \nus\ data reduction as well as the lower energy data from \ch\ and \xmmn, and the multi-wavelength data available for the source; in section \ref{src} we report what was known about NuSTAR J033202--2746.8 from the literature; in section \ref{spec} a detailed X-ray spectral analysis is presented, initially using \nus, \ch\ and \xmmn\ data separately, and subsequently in the broad-band $E\approx0.5-30$ keV energy range, jointly fitting the three datasets. In section \ref{ir} we investigate the IR and radio emission of the source through SED decomposition analysis; discussion and conclusions are in sections \ref{discus} and \ref{concl}, respectively. \begin{figure} \centerline{ \includegraphics[scale=0.5]{fchart_nus_quasar2_JRMimages.jpg}} \caption{\nus\ smoothed images of the source NuSTAR J033202--2746.8 at $3-24$ keV (left), $3-8$ keV (center), and $8-24$ keV (right), from the FPMA and FPMB modules combined. The white circles have 30$''$ radius. The source is undetected down to a probability $Prob=4\times10^{-4}$ in the $3-8$ keV band, but has a clear detection in the broader $3-24$ keV energy band ($Prob=10^{-9}$), and in the $8-24$ keV band ($Prob\approx4\times10^{-6}$).} \label{fig.1} \end{figure} Throughout the paper we assume a cosmological model with $H_0=70\ \rm km\ s^{-1}\ Mpc^{-1}$, $\Omega_M=0.27$ and $\Omega_{\Lambda}=0.73$ \citep{spergel2003}. All the errors are quoted at a 90\% confidence level, unless otherwise specified.	We performed detailed X-ray spectral analysis of NuSTAR J033202--2746.8, the source with the highest band ratio found in the \nus\ observations of the E-CDF-S field so far. The source is very faint in the optical band ($R>25.5$ mag), indicating significant reddening, and bright in the radio band ($L_{\rm 1.4\ GHz}\approx1.2\times10^{27}$ W~Hz$^{-1}$). Using the \nus\ hard X-ray data in combination with existing deep \ch\ and \xmmn\ data, we investigate the broad-band X-ray properties of NuSTAR J033202--2746.8. Moreover, using deep mid- and far-IR data we perform SED decomposition to fully characterize the multi-wavelength properties of the source. Our results can be summarized as follows: \begin{itemize} \item[--] follow-up UV-to-near-IR spectroscopy reveals a faint continuum in the near-IR band with no detection at shorter wavelengths, supporting the idea that the source is obscured in the optical bands. No emission line is detected in the spectrum, preventing a redshift identification from these spectra (however flux losses, and atmospheric absorption might have affected our results). We are planning to perform further follow-up observations, e.g. with the {\it Hubble Space Telescope} (HST) to avoid the atmospheric transmission issues. \smallskip \item[--] Although no secure redshift identification is available from optical/near-IR spectroscopy for the source, we constrain the redshift from the X-ray spectra: $z=2.00\pm0.04$, in agreement with \citet{georgantopoulos2013}. The X-ray luminosity estimated from the X-ray spectra is $L_{\rm 2-10\ keV}\approx10^{44}$ \ergs\ ($L_{\rm 2-10\ keV}\approx4\times10^{44}$ \ergs, corrected for absorption), and $L_{\rm 10-40\ keV}=6.4\times10^{44}$ \ergs, around the peak of the Compton reflection. \smallskip \item[--] From the broad-band X-ray spectral analysis we constrain NuSTAR J033202--2746.8 to be heavily obscured with a column density $N_{\rm H}\approx6\times10^{23}$ cm$^{-2}$, $\sim2-3$ times higher than that previously found using \ch\ or \xmm\ data alone (e.g., \citealt{tozzi2006,castellomor2013,georgantopoulos2013}). \smallskip \item[--] By jointly fitting the \nus, \xmm\ and \ch\ data, and using spectral simulations, we find indications of a Compton reflection component contributing $\sim30$\% to the total emission of NuSTAR J033202--2746.8 at $E\approx10-40$ keV (rest frame), and we estimate the reflection fraction $R=0.55^{+0.44}_{-0.37}$; although this component is relatively weak, it is stronger than that previously found for bright radio-loud quasars, whose X-ray spectra are typically consistent with no reflection.\smallskip % \item[--] The IR SED analysis reveals the mid-IR emission of NuSTAR J033202--2746.8 is dominated by an AGN component, with $\nu L_{\rm 6\ \mu m}\approx3.5\times10^{44}$ \ergs, in agreement with the AGN power estimated in the X-rays, while the far-IR SED is possibly dominated by cool dust emission due to star formation; however, we only place an upper limit on the specific SFR$\lesssim2.1$ Gyr$^{-1}$, which could be consistent with typical star-forming galaxies at $z\approx2$, but also with the lower sSFRs observed in radio-excess and radio-loud AGN. \end{itemize} Although NuSTAR J033202--2746.8 shows some peculiar characteristics, such as the extreme X/O flux ratio (see Sect. \ref{lit}), the lack of optical/near-IR emission lines (Sect. \ref{opt}), and the hard \nus\ band ratio, we conclude that they can be explained through large amount of obscuration around the central black hole or on larger scales. The X-ray spectral properties of NuSTAR J033202--2746.8 are not peculiar, or rare, and could be fairly typical in quasars. We have shown that having higher energy data ($E>10$ keV) is essential to provide a full characterization of the spectral properties of the source, especially at $E\approx20-30$ keV, at the peak of the Compton reflection hump. Such a full spectral characterization is often not feasible when only using lower energy data ($E<10$ keV), but it is essential to make progress on our understanding of the XRB composition and on the AGN population.	14	3	1403.2491	We report Nuclear Spectroscopic Telescope Array (NuSTAR) observations of NuSTAR J033202-2746.8, a heavily obscured, radio-loud quasar detected in the Extended Chandra Deep Field-South, the deepest layer of the NuSTAR extragalactic survey (~400 ks, at its deepest). NuSTAR J033202-2746.8 is reliably detected by NuSTAR only at E > 8 keV and has a very flat spectral slope in the NuSTAR energy band (\Gamma =0.55^{+0.62}_{-0.64}; 3-30 keV). Combining the NuSTAR data with extremely deep observations by Chandra and XMM-Newton (4 Ms and 3 Ms, respectively), we constrain the broad-band X-ray spectrum of NuSTAR J033202-2746.8, indicating that this source is a heavily obscured quasar (N_H=5.6^{+0.9}_{-0.8}\times 10^{23} cm<SUP>-2</SUP>) with luminosity L <SUB>10-40 keV</SUB> ≈ 6.4 × 10<SUP>44</SUP> erg s<SUP>-1</SUP>. Although existing optical and near-infrared (near-IR) data, as well as follow-up spectroscopy with the Keck and VLT telescopes, failed to provide a secure redshift identification for NuSTAR J033202-2746.8, we reliably constrain the redshift z = 2.00 ± 0.04 from the X-ray spectral features (primarily from the iron K edge). The NuSTAR spectrum shows a significant reflection component (R=0.55^{+0.44}_{-0.37}), which was not constrained by previous analyses of Chandra and XMM-Newton data alone. The measured reflection fraction is higher than the R ~ 0 typically observed in bright radio-loud quasars such as NuSTAR J033202-2746.8, which has L <SUB>1.4 GHz</SUB> ≈ 10<SUP>27</SUP> W Hz<SUP>-1</SUP>. Constraining the spectral shape of active galactic nuclei (AGNs), including bright quasars, is very important for understanding the AGN population, and can have a strong impact on the modeling of the X-ray background. Our results show the importance of NuSTAR in investigating the broad-band spectral properties of quasars out to high redshift.	false	[ "NuSTAR J033202", "NuSTAR", "L <SUB>1.4 GHz</SUB", "bright quasars", "high redshift", "the X-ray spectral features", "quasars", "<", "the NuSTAR energy band", "Chandra", "the X-ray background", "the NuSTAR extragalactic survey", "keV", "the NuSTAR data", "bright radio-loud quasars", "the broad-band X-ray spectrum", "The NuSTAR spectrum", "active galactic nuclei", "radio-loud quasar", "Nuclear Spectroscopic Telescope Array (NuSTAR) observations" ]	15.107222	7.032127	119
12398402	[ "Winther, Hans A.", "Ferreira, Pedro G." ]	2015PhRvD..91l3507W	[ "Fast route to nonlinear clustering statistics in modified gravity theories" ]	39	[ "Astrophysics, University of Oxford, DWB, Keble Road, Oxford OX1 3RH, United Kingdom", "Astrophysics, University of Oxford, DWB, Keble Road, Oxford OX1 3RH, United Kingdom" ]	[ "2014JCAP...09..031B", "2015A&A...583A.123G", "2015ApJ...811..116B", "2015JCAP...10..036T", "2015JCAP...12..059B", "2015MNRAS.448.2275A", "2015MNRAS.452.4203M", "2015MNRAS.454.4208W", "2015PhRvD..92f4005W", "2016A&A...595A..78G", "2016JCAP...03..012S", "2016JCAP...11..039L", "2016MNRAS.462.1530A", "2017JCAP...01..047L", "2017JCAP...02..050B", "2017JCAP...03..012V", "2017JCAP...08..006W", "2017PhRvD..95j3515V", "2017PhRvD..96l3526A", "2018IJMPD..2748002L", "2018IJMPD..2748004L", "2018JCAP...12..028V", "2018PhRvD..97b3535V", "2019arXiv190803430B", "2021JCAP...03..010B", "2021JCAP...08..052B", "2021JCAP...09..021F", "2021JCAP...11..050A", "2021PhRvD.103l3525R", "2021RvMP...93a5003B", "2022JCAP...12..028F", "2022LRCA....8....1A", "2022MNRAS.512.1885F", "2023JCAP...03..040W", "2023JCAP...12..045F", "2023arXiv230307121F", "2024A&A...685A.156M", "2024arXiv240510319G", "2024arXiv240613667B" ]	[ "astronomy", "physics" ]	16	[ "98.80.-k", "04.50.Kd", "98.80.Cq", "Cosmology", "Modified theories of gravity", "Particle-theory and field-theory models of the early Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]	[ "1972PhLB...39..393V", "2000PhLB..485..208D", "2002A&A...385..337T", "2002PhRvD..65d4023D", "2003AnPhy.305...96A", "2003JHEP...09..029L", "2004JHEP...06..059N", "2004PhRvD..69d4026K", "2004PhRvD..70l3518B", "2004PhRvL..93q1104K", "2005MNRAS.364.1105S", "2005PhRvD..71d3504W", "2006IJMPD..15.1753C", "2006LRR.....9....3W", "2007PhRvD..75f3501M", "2007PhRvD..76b3514T", "2007PhRvD..76f4004H", "2008PhRvD..78l3523O", "2008PhRvD..78l3524O", "2009IJMPD..18.2147B", "2009PhRvD..79f4036N", "2009PhRvD..79h4003D", "2009PhRvD..80f4023K", "2010PhRvD..82f3519B", "2010PhRvD..82h3503B", "2010PhRvL.104w1301H", "2010arXiv1011.5909K", "2011JHEP...08..108D", "2011PhRvD..83d4007Z", "2012ApJ...748...61D", "2012JCAP...01..051L", "2012JCAP...10..002B", "2012PhLB..715...38B", "2012PhR...513....1C", "2012PhRvD..86d4015B", "2012PhRvD..86l4016B", "2012PhRvL.109x1102K", "2013JCAP...04..029B", "2013JCAP...05..023L", "2013JCAP...06..036T", "2013JCAP...10..027B", "2013JCAP...11..012L", "2013MNRAS.436..348P", "2013PhRvD..88d4057W", "2013arXiv1312.2006K", "2014A&A...562A..78L", "2014PhRvD..89b3521N", "2014PhRvD..90b3507B", "2014PhRvD..90b3508B" ]	[ "10.1103/PhysRevD.91.123507", "10.48550/arXiv.1403.6492" ]	1403	1403.6492_arXiv.txt	The discovery of the accelerated expansion of the Universe is one of the biggest puzzles of modern cosmology and is attributed to an unknown substance dubbed dark energy \cite{2006IJMPD..15.1753C}. One of the proposed solutions to this puzzle is that dark energy is a new field, with a scalar field being the simplest possibility. If such a scalar field exists and has interactions with matter, as is expected from many theories beyond the standard model, then there will be a long-rang fifth-force in nature; i.e., we have a modified theory of gravity \cite{2012PhR...513....1C}. Results from gravity experiments on Earth (see e.g. \cite{2002cls..conf....9A}) and in the solar-system \cite{2006LRR.....9....3W} so far agree perfectly with the predictions of General Relativity (GR) and consequently any modified gravity theory must satisfy the stringent constraints coming from these experiments. This requires either that the scalar field couples to matter much more weakly than gravity or that there exists some mechanism for restoring GR in the solar-system. Over the last decade several different types of screening mechanism have been proposed. The first class of screening is the so-called {\it chameleon} \cite{2004PhRvD..69d4026K,2004PhRvL..93q1104K} and {\it symmetron} \cite{2010PhRvL.104w1301H} mechanism. Here the scalar field is massive and the mass depends on the local matter density. If a body is screened or not depends on the value of its gravitational potential relative to a critical potential defined by the theory. A second class contains models with a shift-symmetry, $\phi \to \phi + c$, and are generally known as {\it k-Mouflage} \cite{2009IJMPD..18.2147B,2011JHEP...08..108D}. In these models screening happens for bodies that experience a large gravitational force (again with respect to a model dependent critical force). A third class contains models with a derivative shift-symmetry, $\partial\phi \to \partial\phi + c$. In this class we find the DGP model \cite{2000PhLB..485..208D,2002PhRvD..65d4023D,2003JHEP...09..029L,2004JHEP...06..059N} (in the decoupling limit) and the Galileon \cite{2009PhRvD..79f4036N,2003AnPhy.305...96A}. Screening in these models takes place for bodies that experience large force-gradients. This is the so-called {\it Vainshtein} mechanism \cite{1972PhLB...39..393V}. We should also mention screening mechanisms for models that employ a disformal coupling to matter \cite{screendisformal}. Here screening is driven by time-derivatives of the scalar field becoming small in high-density environments. The cosmology of modified gravity models, such as the ones mentioned above, has been extensively studied. The main signature they predict, beyond modifying the background cosmology, is to alter structure formation. The scalar fifth-force present in these models is, by design, hidden in high density environments like on Earth and in the solar-system, but in the cosmological background where the density is much smaller the fifth-force can be as strong as gravity leading to potentially large signatures. Because of this effect, to accurately study the effects on structure formation, the first line of attack is linear perturbation theory and naively one would think that on large scales this should be a good approximation. However, linear theory has the disadvantage of not taking the screening mechanism into account and it has been shown (see e.g. \cite{2012JCAP...10..002B,2013JCAP...04..029B}) for many models that linear theory gives a poor fit to the true result, found by solving the full non-linear dynamics in {\it N}-body simulations, even on scales we normally think of as linear. The reason for this is the screening effect on small scales, which represents a breakdown of the super-position principle, making the large-scale fifth-force depend sensitively on small-scale clustering. To calculate accurate predictions for structure formation one is therefore lead to {\it N}-body simulations. Over the last couple of years several codes have been developed to simulate modified gravity models \cite{2013arXiv1307.6748L,2012JCAP...01..051L,2013MNRAS.436..348P,2008PhRvD..78l3523O}. Such simulations need to solve the full Klein-Gordon (KG) equation for the scalar field in order to be able to calculate the fifth-force. Because the KG equation is highly non-linear this task is often hard, in terms of convergence properties, and also computational expensive. A typical {\it N}-body simulation of models in this class can easily take $10$ times as long to finish as a similar simulation for the $\Lambda$CDM model. If modified gravity models such as those discussed above are to be confronted with observations in the non-linear regime then a fast method to compute clustering statistics would be of great value. For $\Lambda$CDM simulations, for example, such a fast method, called COLA, has recently been proposed \cite{2013JCAP...06..036T}. The goal of this paper is to investigate the possibility of a similar speed-up for modified gravity simulations, albeit with a completely different methodology: instead of trying to find a novel way of solving the exact equations we try to construct an equation that can match the behavior in regimes over which we have some analytical control, i.e. in the linear regime and the deep non-linear regime. The field equation we propose is found by combining a screening factor, calculated from spherical symmetric configurations, with the linear Klein-Gordon equation. The screening factor depend only on the metric potential $\Phi$ which is already known to us when performing an {\it N}-body simulation (if we use a particle mesh code) and therefore does not require any additional computations. Importantly, our proposed field equation is linear (in the scalar field) which makes it simple to solve: we can use the same method as used to solve the Poisson equation for $\Phi$. The method proposal therefore has the advantage that it is able to simulate modified gravity theories taking only $\sim 1-2$ times the computational time of a corresponding $\Lambda$CDM simulation\footnote{This estimate is based on our tests on dark matter only simulations, using a particle mesh code, where the computation of the metric potential is the most time consuming part.}. The method is most suitable for particle mesh codes like $\tt{RAMSES}$ \cite{2002AA...385..337T} which we have used in this paper, but in most of our cases it should be fairly straight forward - at least in principle - to implement it in codes that do not calculate the metric potential explicitly, like for example in the popular tree code {\tt{GADGET2}} \cite{2005MNRAS.364.1105S}. The setup of this paper is as follows: in Sec.~\ref{screensec} we briefly review the different screening mechanisms, in Sec.~\ref{appsec} we present our method for the different types of screening mechanisms, then in Sec.~\ref{testsec} we apply the method to {\it N}-body simulations before concluding in Sec.~\ref{concsec}.	\label{concsec} We have proposed a simple and fast, in terms of computational resources needed, method to perform {\it N}-body simulations for scalar-tensor theories which has a screening mechanism on the form described below Eq.~(\ref{lagr}). The method consists of deriving the screening factor from studying at spherical symmetric configurations and rewriting this in terms of the Newtonian potential. This screening factor can then be attached to the linear KG equation and used in simulations. For the three screening mechanisms studied here our method produce accurate results far into the non-linear regime, i.e. up to $k\sim \text{a~few~} h/$Mpc for $f(R)$ gravity and the Galelion and $k\sim 1~h/$Mpc for the symmetron. For the $f(R)$ models we seem to do better the further into the screening-regime we get, i.e. when the linear simulations gets further and further away from the true result. In all test cases our method seems to slightly overestimate the screening (at least on scales $k \lesssim 1~h/$Mpc). The only exception is found for the symmetron. Models where the coupling is field-dependent, such as the symmetron, can have the property that it produces additional screening on small scales. Our method, in its simplest form, does not take this into account and consequently over-estimates the power which is what we find on scales $k\gtrsim 1~h/$Mpc and the mass-function in the high-mass end. Our method can be modified to try to make the fit to the true result better by introducing a fudge-factor that parametrizes this average overestimation. We have only tested our method when it comes to power-spectra and mass-functions. It remains to see how good this method is at predicting other interesting observables such as halo and void profiles, halo and void shapes and velocity statistics to mention some. If attempting to apply our method, another warning is in place: the method is fundamentally phenomenological and if applied it should be tested against full simulations to get an estimate on the error. However, this only needs to be done on a few simulations compared to several tens at least needed to build up a covariance matrix. Our method can also be useful when trying to map out the non-linear regime for a new modified gravity model not simulated before. Using our approach we can very easily implement the model, run simulations, and get a good feel for the possible signatures that it might produce. Finally it will be very interesting to see if our method can be used in conjunction with the COLA approach \cite{2013JCAP...06..036T} to further speed up modified gravity simulations. Such a combined method could open up the window allowing us to do a full MCMC analysis of modified gravity models using data from future large-scale structure surveys.	14	3	1403.6492	We propose a simple and computationally fast method for performing N -body simulations for a large class of modified gravity theories with a screening mechanism such as chameleons, symmetrons, and Galileons. By combining the linear Klein-Gordon equation with a screening factor, calculated from analytical solutions of spherical symmetric configurations, we obtain a modified field equation of which the solution is exact in the linear regime while at the same time taking screening into account on nonlinear scales. The resulting modified field equation remains linear and can be solved just as quickly as the Poisson equation without any of the convergence problems that can arise when solving the full equation. We test our method with N -body simulations and find that it compares remarkably well with full simulations well into the nonlinear regime.	false	[ "nonlinear scales", "full simulations", "screening", "modified gravity theories", "analytical solutions", "spherical symmetric configurations", "Galileons", "N -body simulations", "The resulting modified field equation", "a modified field equation", "linear", "account", "the full equation", "symmetrons", "the nonlinear regime", "the Poisson equation", "chameleons", "the linear regime", "Poisson", "a screening mechanism" ]	10.310741	1.538375	-1
483280	[ "Martin, Nicolas F.", "Ibata, Rodrigo A.", "Rich, R. Michael", "Collins, Michelle L. M.", "Fardal, Mark A.", "Irwin, Michael J.", "Lewis, Geraint F.", "McConnachie, Alan W.", "Babul, Arif", "Bate, Nicholas F.", "Chapman, Scott C.", "Conn, Anthony R.", "Crnojević, Denija", "Ferguson, Annette M. N.", "Mackey, A. Dougal", "Navarro, Julio F.", "Peñarrubia, Jorge", "Tanvir, Nial T.", "Valls-Gabaud, David" ]	2014ApJ...787...19M	[ "The PAndAS Field of Streams: Stellar Structures in the Milky Way Halo toward Andromeda and Triangulum" ]	82	[ "Observatoire astronomique de Strasbourg, Université de Strasbourg, CNRS, UMR 7550, 11 rue de l'Université, F-67000 Strasbourg, France; Max-Planck-Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany;", "Observatoire astronomique de Strasbourg, Université de Strasbourg, CNRS, UMR 7550, 11 rue de l'Université, F-67000 Strasbourg, France", "Department of Physics and Astronomy, University of California, Los Angeles, PAB, 430 Portola Plaza, Los Angeles, CA 90095-1547, USA", "Max-Planck-Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany", "Department of Astronomy, University of Massachusetts, Amherst, MA 01003, USA", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, School of Physics A28, University of Sydney, NSW 2006, Australia", "NRC Herzberg Institute of Astrophysics, 5071 West Saanich Road, Victoria, BC, V9E 2E7, Canada", "Department of Physics and Astronomy, University of Victoria, 3800 Finnerty Road, Victoria, British Columbia V8P 5C2, Canada", "Institute of Astronomy, School of Physics A28, University of Sydney, NSW 2006, Australia", "Department of Physics and Atmospheric Science, Dalhousie University, 6310 Coburg Road, Halifax, NS, B3H 4R2, Canada", "Institute of Astronomy, School of Physics A28, University of Sydney, NSW 2006, Australia", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ, UK; Physics Department, Texas Tech University, Lubbock, TX 79409, USA", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ, UK", "RSAA, The Australian National University, Mount Stromlo Observatory, Cotter Road, Weston Creek ACT 2611, Australia", "Department of Physics and Astronomy, University of Victoria, 3800 Finnerty Road, Victoria, British Columbia V8P 5C2, Canada", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ, UK", "Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK", "Observatoire de Paris, LERMA, 61 Avenue de l'Observatoire, F-75014 Paris, France" ]	[ "2014ApJ...790...74P", "2014ApJ...790L..10G", "2014ApJ...791....9S", "2014ApJ...793...62S", "2014ApJ...795...94B", "2014MNRAS.442.2929V", "2014MNRAS.443L..84B", "2014MNRAS.444.3975D", "2015A&A...584A.120B", "2015ApJ...801..105X", "2015ApJ...802L..18L", "2015ApJ...803...80K", "2015ApJ...814L...7K", "2015MNRAS.448..855Y", "2015MNRAS.452..676P", "2015MNRAS.454.3613S", "2016A&A...589A.120K", "2016A&A...592A..55V", "2016ASSL..420...63Y", "2016ASSL..420...87G", "2016ApJ...816...80J", "2016ApJ...818...39H", "2016ApJ...818...40M", "2016ApJ...818...41S", "2016ApJ...818..194N", "2016ApJ...819..124L", "2016ApJ...820...18C", "2016ApJ...824...35T", "2016MNRAS.456..405M", "2016MNRAS.460L.114M", "2016MNRAS.461.1590E", "2016MNRAS.461.4374M", "2016MNRAS.463..102E", "2016MNRAS.463.1759B", "2017A&A...601A..37B", "2017ApJ...840L..25M", "2017ApJ...848..128I", "2017Galax...5...44J", "2017JCAP...11..003B", "2017MNRAS.467.2636D", "2018AJ....156..230C", "2018ApJ...853...29K", "2018ApJ...867..144M", "2018ApJ...868...55M", "2018MNRAS.473..647D", "2018MNRAS.473.1461P", "2018MNRAS.477.2419M", "2018MNRAS.478.3862M", "2019A&A...621A..13V", "2019A&A...632L..13M", "2019ApJ...884...67W", "2019ApJ...886..113S", "2019MNRAS.482.2795H", "2019MNRAS.484.1756M", "2019MNRAS.486..523B", "2019MNRAS.486..823R", "2019MNRAS.490.2905P", "2020A&A...644A..42R", "2020ApJ...893..120Z", "2020MNRAS.491.5101A", "2020MNRAS.494..983R", "2020SCPMA..6309801W", "2021A&A...654A..23G", "2021ApJ...923..146G", "2021MNRAS.502.4170R", "2021MNRAS.503..472O", "2021MNRAS.504.3098P", "2021MNRAS.507.2592S", "2022A&A...660A..29T", "2022A&A...662A.124S", "2022A&A...667A..37Y", "2022ApJ...930..103Y", "2022ApJ...941...19P", "2022MNRAS.510L..13L", "2022MNRAS.516.4469Z", "2022PASP..134h4101B", "2023A&A...677A..91K", "2023ApJ...949...48A", "2024ApJ...962..151A", "2024MNRAS.530...95G", "2024arXiv240314234O", "2024arXiv240519410B" ]	[ "astronomy" ]	8	[ "Galaxy: halo", "Galaxy: structure", "Local Group", "Astrophysics - Astrophysics of Galaxies" ]	[ "2001ApJ...559..716T", "2002AJ....124.1452F", "2002MNRAS.333..779P", "2003MNRAS.340L..21I", "2004ApJ...615..732R", "2004ApJ...615..738M", "2005ApJ...626..128P", "2005ApJ...634..287I", "2006ApJ...641L..37G", "2006ApJ...642L.137B", "2006ApJ...650L..33P", "2006ApJ...653..255C", "2007AJ....134..376D", "2007ApJ...668L.123M", "2007ApJ...671.1591I", "2007ApJS..172..545R", "2007MNRAS.380..281M", "2008AJ....135.1998R", "2008ApJ...680..295B", "2008ApJ...684.1075M", "2008ApJ...688..254K", "2008ApJ...688.1009M", "2008ApJ...689..184M", "2008ApJ...689..936J", "2009ApJ...693.1118G", "2009MNRAS.397L..26C", "2009Natur.461...66M", "2010AJ....140..962M", "2010ApJ...714..663D", "2010ApJ...714L..12M", "2010MNRAS.406..744C", "2011ApJ...730L...6S", "2011ApJ...733L...7H", "2011ApJ...738...98G", "2011Natur.477..301P", "2012ApJ...750L..41W", "2012ApJ...752...45T", "2012ApJ...754..101C", "2012ApJ...760L...6B", "2012MNRAS.425.2335S", "2012MNRAS.427..127B", "2013ApJ...765L..39M", "2013ApJ...768..172C", "2013ApJ...776...80M", "2013MNRAS.429..159G", "2013MNRAS.435.1928P", "2013MNRAS.436.3231K", "2014A&A...564A..18P", "2014ApJ...780..128I", "2014MNRAS.440.2564F" ]	[ "10.1088/0004-637X/787/1/19", "10.48550/arXiv.1403.4945" ]	1403	1403.4945_arXiv.txt	The mapping of the structured nature of the stellar halo of the Milky Way (MW) is one of the greatest successes of the Sloan Digital Sky Survey (SDSS). In their `Field of Streams' of the northern Galactic cap, \citet{belokurov06b} highlighted at least four distinct stellar structures that are mainly thought to be the result of the accretion of dwarf galaxies onto the Milky Way over the last few billion years \citep{bell08}. This process of hierarchical formation is a tenet of the currently favored model for the formation of galactic outskirts and is also observed in other nearby galaxies that we can study with enough detail \citep[e.g.][]{ferguson02,martinez-delgado08,martinez-delgado10,mouhcine10}. For instance, the Pan-Andromeda Archaelogical Survey (PAndAS; \citealt{mcconnachie09}) covers a large fraction of the stellar halo of the Andromeda galaxy (M31), and also reveals a wealth of stellar features produced by the tidal disruption of dwarf galaxies \citep{ibata14}. Although it requires covering large swaths of the sky, imaging the outskirts of the Milky Way remains beneficial to find and track the faintest of stellar streams that cannot be discovered beyond a few tens of kiloparsecs. Discovering these, or proving their absence, is essential if we want to robustly compare the properties of modelled stellar halos with reality \citep{johnston08,cooper10}. Further analyses of the SDSS have illustrated that the Field of Streams is not the whole story, with the discovery of additional thin and/or faint streams from main sequence (MS) and main sequence turn-off (MSTO) stars \citep[e.g.][]{grillmair06,grillmair09,grillmair11,bonaca12,martinc13}. Yet the depth limits of the SDSS and similar surveys like Pan-STARRS~1 are quickly reached in such studies. Until the era of the Large Synoptic Survey Telescope (LSST), almost a decade from now, targeted but deeper surveys can be advantageous for the mapping of the MW stellar halo \citep[e.g.][]{robin07,pila-diez14}, although interpretation can be difficult if angular coverage is limited. Of particular interest is the region towards M31, which is a target of interest for studies of our cosmic neighbor. These surveys have yielded detections of the Monoceros Ring \citep[MRi;][]{ibata03}, a stellar structure at a heliocentric distance of $\sim7\kpc$ that appears to be circling the disk in the second and third Galactic quadrants and whose nature is a matter of debate (see \citealt{conn12}, and references therein). Beyond the MRi, \citet{majewski04b} discovered a fainter and farther stellar feature, confirmed via the spectroscopy of 2MASS red giant branch stars \citep{rocha-pinto04}. This structure, Triangulum-Andromeda (TriAnd), is estimated to be located at a heliocentric distance of 16 to $25\kpc$, and shows a cloud-like morphology that covers hundreds of square-degrees. This was later confirmed via the analysis of the MW stellar halo in the CFHT pilot program that would become PAndAS \citep{martin07b}. At that time, the imaging data, which only covered a contiguous $\sim76\textrm{ deg}^2$ ($\sim20$\%) of the final survey footprint, showed no spatially well-defined stellar structure in the Milky Way halo, but allowed for the detection of yet another structure, Triangulum-Andromeda 2 (TriAnd2), as an even more distant MS feature in the color-magnitude diagram (CMD). Neither TriAnd, nor TriAnd2 located at a heliocentric distance of $\sim28\kpc$, showed strong spatial variations within the coverage of the survey. In this paper, we revisit our mapping of the stellar halo of the MW in the region toward M31. We use a large fraction of the full PAndAS coverage ($\sim360\textrm{ deg}^2$) to trace the structure of halo MS and MSTO stars. In Section~2, we briefly describe the PAndAS data we use, while Section~3 presents our mapping of the MW halo and the global properties of the structures. Section~4 focusses on a well-defined stellar stream for which we highlight a possible progenitor and a serendipitous velocity measurement. Finally, we conclude in Section~5. Where necessary, the distance to the Galactic center is assumed to be $8\kpc$.	From the analysis of blue MS and MSTO stars in the PAndAS footprint, we have shown that the Milky Way halo is very structured in the direction of Andromeda and Triangulum, out to distances of at least $\sim30\kpc$. In addition to the stellar features already known in this part of the sky --- the Monoceros Ring \citep{ibata03}, Triangulum/Andromeda \citep{majewski04b,rocha-pinto04}, Triangulum/Andromeda~2 \citep{martin07b}, the Pisces-Triangulum globular cluster stream \citep{bonaca12,martinc13} --- we reveal the presence of at least two new stellar features. At $D_\sun\sim17\kpc$, a stream-like structure crosses the PAndAS footprint, mainly along the east-west axis that is coplanar with the Galactic disk, and another `blob,' which is cut off by the north-eastern limits of the survey. The stellar populations that compose these two structures are very similar, in both cases old and metal-poor, but there is no obvious connection between them. The stream exhibits a strong radial distance gradient, with its heliocentric distance diminishing by $2\kpc$ over a sky-projected path of only $\sim3\kpc$. It has a width of a few hundred parsecs and a depth smaller than $1.8\kpc$, implying it is likely the product of the disruption of a dwarf galaxy, whose remnant may be a localized overdensity embedded in the stream, located at $(\ell,b)\sim(117.2\deg,-16.6\deg)$. We further find a kinematic signature of likely stream stars in our library of DEIMOS/Keck spectra targeting M31 stellar structure. We show that the PAndAS MW stream star candidates are dynamically cold with a velocity dispersion of no more than $7.1\kms$ at the 90-percent confidence level. This analysis of the PAndAS foreground reveals that the SDSS `Field of Streams' \citep{belokurov06b} is but the tip of the iceberg. Reaching fainter/more diffuse structures, even over the comparatively small PAndAS area ($\sim360\textrm{deg}^2$), unveils many substructures which would be difficult to map convincingly without the exquisite PAndAS photometry. Yet PAndAS does not even cover 1 percent of the night sky. One can therefore expect that surveys like the DES\footnote{http://www.darkenergysurvey.org}, or the LSST\footnote{http://www.lsst.org/lsst/}, will reveal another wealth of stellar streams within at least $50\kpc$ from the Galactic center. This is not unexpected from simulations of the hierarchical build-up of stellar halos \citep[e.g.][]{johnston08,helmi11}, and neither are the morphological differences between the structures. In the PAndAS Field of Streams, we find stream-like, cloud-like, and `blob'-like features, as one can expect at these distances that transition between the very inner parts of the halo, for which dynamical times are short, and its outer regions, for which these can reach a Hubble time. Finally, the discovery of the new stream is testament that, even though most streams known to date are on orbits close to polar \citep{pawlowski13a}, more planar streams do exist. Our tally of streams may well be biased away from planar orbits, especially if dynamical friction is efficient at aligning the orbit of non-polar accretions with the MW plane \citep{taylor01,penarrubia02,penarrubia06}, effectively hiding these accretions from most surveys that shy away from the disk-dominated regions.	14	3	1403.4945	We reveal the highly structured nature of the Milky Way (MW) stellar halo within the footprint of the Pan-Andromeda Archaeological Survey (PAndAS) photometric survey from blue main sequence (MS) and MS turn-off stars. We map no fewer than five stellar structures within a heliocentric range of ~5-30 kpc. Some of these are known (the Monoceros Ring, the Pisces/Triangulum globular cluster stream), but we also uncover three well-defined stellar structures that could be, at least partly, responsible for the so-called Triangulum/Andromeda and Triangulum/Andromeda 2 features. In particular, we trace a new faint stellar stream located at a heliocentric distance of ~17 kpc. With a surface brightness of Σ<SUB> V </SUB> ~ 32-32.5 mag arcsec<SUP>-2</SUP>, it follows an orbit that is almost parallel to the Galactic plane north of M31 and has so far eluded surveys of the MW halo as these tend to steer away from regions dominated by the Galactic disk. Investigating our follow-up spectroscopic observations of PAndAS, we serendipitously uncover a radial velocity signature from stars that have colors and magnitudes compatible with the stream. From the velocity of eight likely member stars, we show that this stellar structure is dynamically cold, with an unresolved velocity dispersion that is lower than 7.1 km s<SUP>-1</SUP> at the 90% confidence level. Along with the width of the stream (300-650 pc), its dynamics point to a dwarf-galaxy-accretion origin. The numerous stellar structures we can map in the MW stellar halo between 5 and 30 kpc and their varying morphology is a testament to the complex nature of the stellar halo at these intermediate distances.	false	[ "stars", "the MW stellar halo", "a new faint stellar stream", "blue main sequence", "MW", "the stellar halo", "MS", "The numerous stellar structures", "this stellar structure", "Galactic", "the Pan-Andromeda Archaeological Survey", "PAndAS", "the MW halo", "the Milky Way (MW) stellar halo", "magnitudes", "the complex nature", "colors", "M31", "their varying morphology", "Triangulum/Andromeda" ]	9.636232	7.543697	-1
678287	[ "Villa, Eleonora", "Matarrese, Sabino", "Maino, Davide" ]	2014JCAP...06..041V	[ "Cosmological dynamics: from the Eulerian to the Lagrangian frame. Part I. Newtonian approximation" ]	12	[ "Dipartimento di Fisica, Università di Milano, via Celoria 16, 20154 Milano, Italy", "Dipartimento di Fisica e Astronomia ``G. Galilei\", Università degli Studi di Padova and INFN Sezione di Padova, via Marzolo 8, 35131 Padova, Italy; Gran Sasso Science Institute, INFN, viale F. Crispi 7, 67100 L'Aquila, Italy;", "Dipartimento di Fisica, Università di Milano, via Celoria 16, 20154 Milano, Italy" ]	[ "2014CQGra..31w4004R", "2014CQGra..31w4005V", "2014PhRvD..90l3503R", "2015EPJC...75..381E", "2015JCAP...07..036E", "2016JCAP...01..030V", "2016PDU....14...11B", "2016PhRvD..94h3515R", "2018PhRvD..98d3507L", "2020PDU....2900565C", "2022JCAP...03..048E", "2024arXiv240503397N" ]	[ "astronomy" ]	9	[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1971grc..conf..104E", "1978PhRvD..17.2529S", "1980PhRvD..22.1882B", "1989A&A...223....9B", "1992ApJ...394L...5B", "1993PhRvD..47.1311M", "1994ApJ...435....1B", "1994MNRAS.271..513M", "1994PhRvL..72..320M", "1995ApJ...442...30K", "1995ApJ...445..958B", "1995MNRAS.276...39C", "1995MNRAS.276..115C", "1996MNRAS.283..400M", "1997CQGra..14.2585B", "1998PhRvD..58d3504M", "2005JCAP...10..010B", "2005PhRvD..71b3524K", "2005PhRvD..71d3508C", "2009ApJ...706L..91V", "2009PhRvD..79l3507W", "2009PhRvD..80h3514Y", "2010CQGra..27l4009B", "2011JCAP...04..011B", "2011JCAP...08..024V", "2011JCAP...10..031B", "2011PhRvD..83l3505C", "2011PhRvD..83l3519M", "2011PhRvD..84f3505B", "2012JCAP...07..051B", "2012JCAP...10..025B", "2012JCAP...11..045B", "2012MNRAS.419.1937M", "2012PhRvD..85b3504J", "2012PhRvD..85d1301B", "2012PhRvD..86b3520B", "2013JCAP...02..044M", "2013JCAP...06..002B", "2013JCAP...11..019F", "2013PhRvD..87l3525R", "2013PhRvD..88j3527A", "2013PhRvL.110b1301B", "2014CQGra..31t5001U", "2014JCAP...05..022P", "2014PhRvD..89d4010B", "2014PhRvD..89f3509R", "2014PhRvD..89f3543A", "2015JCAP...09..021T", "2015MNRAS.446..677R" ]	[ "10.1088/1475-7516/2014/06/041", "10.48550/arXiv.1403.6806" ]	1403	1403.6806_arXiv.txt	It is believed that the distribution of matter in the early Universe was very smooth, the best indication being the tiny fluctuations in the Cosmic Microwave Background. However, the distribution of matter in the Universe at present time is inhomogeneous on scales below about 100 $h^{-1}$ Mpc ($h$ being the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$) and the gravitational dynamics is non-linear below about 10 $h^{-1}$ Mpc. The theoretical study of structure formation connects the early Universe with that observed today. The gravitational instability is governed by the equations provided by General Relativity (GR) but, for most purposes, we can use the Newtonian approximation, namely the weak-field and slow-motion limit of GR. These conditions are verified at small scales, well inside the Hubble radius, where the peculiar gravitational potential $\varphi_g$, divided by the square of the speed of light to obtain a dimensionless quantity, remains much less than unity, while the peculiar velocity never become relativistic. The peculiar gravitational potential is related to the matter-density fluctuation $\delta$ via the cosmological Poisson equation \begin{equation} \nabla^2\varphi_g=4\pi G a^2 \rho_b \delta \end{equation} where $\rho_b$ is the Freidmann-Robertson-Walker (FRW) background matter density, $\delta=\left(\rho-\rho_b\right)/\rho_b$ is the density contrast and $a(t)$ is the scale-factor, which obeys the Friedmann equation. This implies that for a fluctuation of proper scale $\lambda$, \begin{equation} \frac{\varphi_g}{c^2} \sim \delta\left(\frac{\lambda}{r_H}\right)^2 \;, \end{equation} where $r_H=cH^{-1}$ is the Hubble radius. This very fact tells us that the weak-field approximation does not necessarily imply small density fluctuations, rather it depends on the ratio of the perturbation scale $\lambda$ to the Hubble radius. That is why Newtonian gravity is widely used to study structure formation at small scales, also in the non-linear regime. The evolution of perturbations is dealt with analytically, within perturbation theory, or numerically, by means of N-body simulations. However, there are situations, even within the Hubble horizon, where the Newtonian treatment is not well suited. The Newtonian approximation of GR, which consists in perturbing the time-time component of the space-time metric by an amount $2\varphi_g/c^2$ fails to produce an accurate description of photon trajectories: it is well know that the Newtonian estimate of the Rees-Sciama effect and of gravitational lensing is incorrect by a factor of 2. The correct calculation involves the weak-field limit of GR, which is valid for slow motions of the sources of the gravitational field, but allows test particles to be relativistic. The related metric is perturbed also in the space-space component of the space-time metric by an amount $-2\varphi_g/c^2$. Above all, the upcoming galaxy surveys will probe increasingly large scales, approaching the Hubble radius, where the Newtonian approximation is no longer valid and a relativistic treatment is therefore mandatory, both to study the grow of structures and to analyse cosmological observations at these scales. The study of gravitational instability in the non-linear regime and in GR is of course very challenging and it cannot be accomplished without some approximations. The most widely used is cosmological relativistic perturbation theory, truncated at first or at second-order. In the perturbative framework, the second-order approach in calculations is crucial: the linear modes generate non-linear ones, many of which are not taken into account in Newtonian theory, and play an important role in the dynamics of perturbations as well as in the calculations of the gravitational lensing and redshift space distortions, see ref.~\cite{veneziano1}-\cite{umeh} and refs. therein. On the other hand, a description of the fully non-linear dynamics of the perturbations can be carried out by means of effective fluid description for small-scale non-linearities in the framework of GR, as proposed in ref.~\cite{baumann} or by means of an effective field theory approach, as proposed in ref.~\cite{porto} using a Lagrangian approach. Also, exact analytical solutions of Einstein's equations, can be important tools to investigate non-linear effects, e.g. in the path of photons throughout inhomogeneities, although in simplified contexts or in presence of some symmetry in the problem, see e.g. refs.~\cite{meures1} and~\cite{meures2}. Recently, a GR analysis of galaxy clustering, refs.~\cite{BMR2005}-\cite{raccanelli} and refs. therein, has drawn attention to other relativistic effects in the observations, such as the gauge dependence of the galaxy bias, which become sizeable near the Hubble horizon, see ref.~\cite{brunibias} and ref.~\cite{wands}. In this respect, a second-order calculation actually shows a gauge-dependence in the matter density also in the Newtonian expressions, see ref.~\cite{BMPR2010} and ref.~\cite{colombi92}. The connection between Newtonian N-body simulations and GR is also of great interest today: in ref.~\cite{chisari} the authors provide a dictionary for how to interpret the outputs of numerical simulation, run using Newtonian dynamics, with respect to GR at first order. In ref.~\cite{durrerNbody} a formalism for GR N-body simulations which go beyond standard perturbation theory is provided in the weak-field limit of GR, which is then used in ref.~\cite{durrerdlz} to compute the distance-redshift relation for a plane-symmetric universe. For the interpretation of N-Body simulations and the connection to the Lagrangian and Eulerian perspectives see ref.~\cite{RigopoulosNbody}, where the relativistic corrections to Newtonian cosmology are analysed with the gradient expansion method on scales close to the Hubble scale.\\ A fully second-order implementations of GR corrections to N-body simulations is still missing and would certainly be useful, although second-order perturbation theory should not to be considered as an exhaustive study for non-linear dynamics. Rather, we should seek for a relativistic and non-perturbative approach, capable to disentangle the Newtonian from the GR contributions. An alternative approximation scheme is well-suited for this purpose: the Newtonian analysis can be improved with the post-Newtonian (PN) approximation of GR, which provides the first relativistic corrections for a system of slowly moving particles bound together by gravitational forces and thus it can be used to account for the moderately non-linear gravitational field generated during the highly non-linear stage of the evolution of matter fluctuations on intermediate scales. It is a crucial improvement of both the aforementioned approximations, as it could bridge the gap between relativistic perturbation theory and Newtonian structure formation, providing a unified approximation scheme able to describe the evolution of cosmic inhomogeneities from the largest observable scales to small ones, including also the intermediate range, where the relativistic effects cannot be ignored and non-linearity starts to be relevant. Nonetheless, few attempts have been made so far to go beyond the Newtonian approximation on non-linear scales. A relevant difficulty of this scheme is that in the PN framework the background is not merely the FRW metric but a Newtonian metric yet describing non-linearities. This very fact and has so far prevented from proceeding in this direction because of its computational complexity, except for symmetric situations. In ref.~\cite{VMM} the PN solution of the Einstein field equations describing the non-linear cosmological dynamics of plane-parallel perturbations in the synchronous and comoving gauge was obtained. It extends the Zel'dovich approximation, which, in turn, is exact for non-linear plane-parallel dynamics in Newtonian gravity. The PN approximation has by construction a direct correspondence with Newtonian quantities: the PN expressions are sourced only by the non-linear Newtonian terms which can be extracted e.g. from N-body simulations (or by approximate analytical expressions obtained via the Zel'dovich approximation). For a recent attempt to include PN-type corrections from N-body simulations see ref.~\cite{brunidrag} and ref.~\cite{brunilensing}. In this paper we consider the non-linear dynamics of cosmological perturbations of an irrotational collisionless fluid, the FRW background being the Einstein-de Sitter model. We discuss the connection between GR and Newtonian gravity in the Eulerian and in the Lagrangian picture. Then we provide the transformation rule between the Eulerian and the Lagrangian frames: it is fully four-dimensional and non-perturbative and clarifies the role of the transformation of the time and spatial coordinates. Our approach here is different from the standard perturbative gauge transformation in GR, see refs.~\cite{MMB}, \cite{MBgeo} and \cite{bardeen}, and from the spatial coordinate transformation of Newtonian theory. Most derivations of the Newtonian limit of GR are coordinate-dependent, thus a precise understanding of the Newtonian correspondence between the Eulerian and the Lagrangian frame has to be considered as the starting point e.g. for studying the gauge dependence when we want to add GR corrections in a perturbed space-time from a non-perturbative perspective. In a forthcoming paper we will extend this analysis to the PN approximation. This paper is organised as follows: in the next two sections we review the study of the non-linear dynamics in the Newtonian limit of GR in the Eulerian and Lagrangian pictures, respectively. In section 4 we present the coordinate transformation and in section 5 the second-order expansion of our quantities in perturbation theory. We first review the Newtonian results in the Eulerian picture. Then we show that our calculation fully recovers the solution for the Lagrangian metric in the Newtonian limit. Conclusions are drawn in the last section. Indices notation: we use $A, B, ...$ for spatial Eulerian indices, $\a, \b, ...$ for spatial Lagrangian and $a, b, ...$ to indicate space-time indices in any gauge. The superscript or subscript $\mathcal{E}$ and $\mathcal{L}$ stands for Eulerian and Lagrangian, respectively, when needed for clarity.	ds^2=a^2\left[-\left(1+2\frac{\varphi_g^\mathcal{E}}{c^2}\right)c^2 d\eta^2 +\delta_{AB}dx^A dx^B\right]. \end{equation} We then transformed this metric to the Lagrangian frame and obtained the Newtonian line-element in the synchronous and comoving gauge as in ref.~\cite{mater}, which is given by \begin{equation}\label{gLconcl} ds^2=a^2\left[-c^2 d\tau^2 +\d_{AB}\mathcal{J}^{A}_{\a}\mathcal{J}^{B}_{\b} dq^\a dq^\b\right]\,. \end{equation} As we said, our starting point was the metric of eq. \eqref{gEconcl} in the Poisson gauge, which we dubbed Newtonian, as the metric variables appearing there are just those needed for the Newtonian equations of motion. With our transformation we arrived at a Newtonian metric in the synchronous and comoving gauge, where, once again, we have just the variables needed for the Newtonian Lagrangian equations of motion, namely the spatial Jacobian matrix. However, the Newtonian three-dimensional space has vanishing spatial curvature, therefore, the Newtonian Lagrangian metric, eq. \eqref{gLconcl}, can be transformed globally to the Einstein-de Sitter background metric, the transformation being just the spatial transformation to Eulerian coordinates $dx^A=\mathcal{J}^a_\a dq^\a$, without changing the time coordinate. As we will see below, this inconsistency can be easily overcome by requiring that the Ricci four-dimensional curvature scalar is preserved by the transformation, as it should. As we have shown, the Lagrangian frame can be transformed to the locally flat inertial frame by means of the transformation to coordinates $\left(\tau+1/c^2\,\xi^0, x^A\right)$: \begin{equation} ds^2=a^2\left[-c^2 \left(d\tau+\frac{1}{c^2}d\xi^0\right)^2+\d_{AB}\mathcal{J}^{A}_{\a}\mathcal{J}^{B}_{\b} dq^\a dq^\b\right]\,. \end{equation} When our solution $\xi^0$, \begin{equation} \label{xi0solutionconcl} \xi^0_\mathcal{L}=\frac{1}{a}\int^\eta_{\eta_{in}}a\left(-\varphi_g^\mathcal{L}+\frac{1}{2}v^Av^B\delta_{AB}\right) d\tilde{\eta} +\frac{C(q^\a)}{a}\, \end{equation} with $C(q^\a)$ fixed from \begin{equation} v^{K}=\frac{\partial \xi^0_\mathcal{L}}{\partial q^\lambda}\mathcal{J}_F^\lambda\delta^{FK}\label{dxxi0concl}\,, \end{equation} is expanded in $1/c^2$, at lowest order this line-element reproduces the Newtonian one in the Eulerian frame, eq. \eqref{gEconcl}. On the other hand, the spatial scalar curvature $^{(3)}\!{\cal R}$ vanishes at lowest order in both frames. The conformal four-dimensional scalar curvature\footnote{We only give the conformal curvature here because, apart from the $a^{-2}$ factor, the extra term in the physical curvature is simply the Einstein-de Sitter scalar curvature in both gauges}, $^{(4)}\!{\cal R}\equiv\mathcal{R}$ of the metric \eqref{gEconcl} is given by \begin{equation} \mathcal{R}=-2\partial^A\partial_A\varphi_g^\mathcal{E} \end{equation} On the other hand, the conformal four-dimensional curvature in the synchronous and comoving gauge is given by \begin{equation} \label{R4Dsynchro} \mathcal{R}=2\vartheta'+\vartheta^2+\vartheta^\mu_\nu\vartheta^\nu_\mu+c^2\,^{(3)}\!{\cal R} \end{equation} and, at lowest order in our $1/c^2$ expansion, the PN spatial curvature contributes to the four-dimensional $\mathcal{R}$ \begin{equation} \mathcal{R}=2\overline{\vartheta}'+\overline{\vartheta}^2+\overline{\vartheta}^\mu_\nu\overline{\vartheta}^\nu_\mu+\,^{(3)}\!{\cal R}^{PN}\,. \end{equation} In other words, at lowest order in our $1/c^2$ expansion, in the Eulerian frame, only the perturbation of the time-time component of the metric contributes to the scalar curvature $\mathcal{R}$, whereas in the Lagrangian frame we need also the spatial PN term coming from the spatial PN metric: setting $\,^{(3)}\!\mathcal{R}=0$ in eq. \eqref{R4Dsynchro} at the lowest order would be incorrect. This very fact is clear from the PN expansion which actually shows that the metric contributes at different orders to the four-dimensional curvature.\\ In order to find the correct four-dimensional curvature in the Lagrangian frame, one cannot use the metric in eq. \eqref{gLconcl}, since the required PN part is missing. This means that, in order to obtain the same expression for the scalar curvature $\mathcal{R}$, at lowest order after the change of frame, we need to start from the weak-field metric in the Poisson gauge: \begin{equation} ds^2=a^2\left[-\left(1+2\frac{\varphi_g^\mathcal{E}}{c^2}\right)c^2 d\eta^2 +\left(1-2\frac{\varphi_g^\mathcal{E}}{c^2}\right)\delta_{AB}dx^A dx^B\right]\,, \end{equation} where only the scalar PN mode is considered in the metric, since vector and tensor modes give higher-order contributions to $\mathcal{R}$, once transformed to the Lagrangian frame. The transformation to the Lagrangian frame finally leads to \begin{equation} ds^2=a^2\left[-c^2 d\tau^2 +\left(1+\frac{\chi}{c^2}\right)\mathcal{J}^{A}_{\a}\mathcal{J}^{B}_{\b}\d_{AB}dq^\a dq^\b\right], \end{equation} where, see ref.~\cite{mater}, \begin{equation} \chi= 2\mathcal{H}\xi^0_\mathcal{L} -2\varphi_g^\mathcal{L}-\Upsilon^\mathcal{L}. \end{equation} In the latter expression, the potential $\Upsilon^\mathcal{L}$ is given by \begin{equation} \overline{\mathcal{D}}^\ss\overline{\mathcal{D}}_\ss\Upsilon^\mathcal{L}=-\frac{1}{2}\left(\overline{\vartheta}^2-\overline{\vartheta}^\mu_\nu\overline{\vartheta}_\mu^\nu\right), \end{equation} and the peculiar velocity-gradient tensor is written in terms of our solution $\xi^0_\mathcal{L}$, eq. \eqref{xi0solutionconcl}, as $\overline{\vartheta}^\mu_\nu=\overline{\mathcal{D}}^\mu\overline{\mathcal{D}}_\nu\xi^0_\mathcal{L}$. The PN scalar mode $\chi$ comes from the transformation of the time coordinate, keeping only the scalar contributions in the PN spatial metric. Following this procedure, starting from the scalar curvature $\mathcal{R}=-2\partial^A\partial_A\varphi_g^\mathcal{E}$ in the Eulerian frame we arrive at the same expression $\mathcal{R}=-2\overline{\mathcal{D}}^\a\overline{\mathcal{D}}_\a\varphi_g^\mathcal{L}$ in the Lagrangian frame. Let us now compare our approach with that of ref.~\cite{MPS1993} and ref.~\cite{MPS1994}. Both papers deal with relativistic dynamics and consider the parallel transport condition, eq. \eqref{paralleltrans}, which in the Lagrangian approach of the synchronous and comoving gauge becomes \begin{equation} \mathcal{J}^{A'}_\a=\vartheta^\ss_\a\mathcal{J}^{A}_\ss\,. \end{equation} Then they solve the Einstein equations and find the velocity-gradient tensor\footnote{In ref.~\cite{MPS1993} the Einstein equations are integrated numerically in the special case of the so-called "silent Universe", see refs. \cite{MPSprl94} and \cite{BMP95}, and solved analytically for the plane-parallel dynamics. In ref.~\cite{MPS1994} the calculation is performed at second order in relativistic perturbation theory.} and the Jacobian matrix. Finally the Eulerian trajectories and peculiar velocity are obtained from \begin{equation} dx^A= \mathcal{J}^{A}_\ss dq^\ss\,. \end{equation} In this paper instead we have reconstructed the Lagrangian dynamics from the Eulerian fields in the Newtonian limit, thus our procedure actually goes in the opposite direction. To conclude, let us emphasise once again that our transformation is different from the standard perturbation theory one. In standard perturbation theory, the background spatial transformation is simply $dx^A=\delta^A_\a dq^\a$; in a relativistic calculation, the same background transformation from the Eulerian to Lagrangian frame is simply the standard Lorentz transformation, with boost velocity equal to the first-order peculiar velocity, see ref.~\cite{bert&hamilton}. Instead, in the Newtonian limit, we have performed a non-trivial spatial transformation from the Eulerian to the Lagrangian frame -- the two frames moving with relative velocity $v^A$ -- and we have subsequently changed the parametrisation of the hyper-surfaces from the rest frame of Eulerian observers to the rest frame of the matter. This transformation is fully non-perturbative and will provide the background transformation for the same procedure in the PN approach. A different approach to the Newtonian limit, see ref.~\cite{brunidrag}, is to consider the lower order in the $1/c^n$ expansion of all the Einstein equations, thus including in the Newtonian metric the scalar $\mathcal{O}(1/c^2)$ in the spatial part and the vector $\mathcal{O}(1/c^3)$ in the time-space component in the Poisson gauge. The difference between these two approaches is just semantic: in the Poisson gauge, what we call PN in our viewpoint is the same as what is called Newtonian in ref.~\cite{brunidrag}. The present transformation has to be considered the background transformation for the PN expansion. Nevertheless, it clarifies the meaning of the change of space and time coordinates from the Eulerian to the Lagrangian frame.	14	3	1403.6806	We analyse the non-linear gravitational dynamics of a pressure-less fluid in the Newtonian limit of General Relativity in both the Eulerian and Lagrangian pictures. Starting from the Newtonian metric in the Poisson gauge, we transform to the synchronous and comoving gauge and obtain the Lagrangian metric within the Newtonian approximation. Our approach is fully non-perturbative, which implies that if our quantities are expanded according to the rules of standard perturbation theory, all terms are exactly recovered at any order in perturbation theory, only provided they are Newtonian. We explicitly show this result up to second order and in both gauges. Our transformation clarifies the meaning of the change of spatial and time coordinates from the Eulerian to the Lagrangian frame in the Newtonian approximation.	false	[ "Newtonian", "Lagrangian", "standard perturbation theory", "perturbation theory", "General Relativity", "the Newtonian approximation", "the Newtonian metric", "second order", "the Newtonian limit", "the Lagrangian metric", "Eulerian", "Poisson", "the Lagrangian frame", "the Poisson gauge", "the synchronous and comoving gauge", "both gauges", "both the Eulerian and Lagrangian pictures", "the non-linear gravitational dynamics", "spatial and time coordinates", "second" ]	10.570685	1.372687	-1
813631	[ "Capozziello, Salvatore", "Farooq, Omer", "Luongo, Orlando", "Ratra, Bharat" ]	2014PhRvD..90d4016C	[ "Cosmographic bounds on the cosmological deceleration-acceleration transition redshift in f(R) gravity" ]	133	[ "Dipartimento di Fisica, Università di Napoli \"Federico II\", Via Cinthia, I-80126 Napoli, Italy; INFN Sezione di Napoli, Monte Sant'Angelo, Via Cinthia, I-80126 Napoli, Italy; Gran Sasso Science Institute (INFN), Viale Francesco Crispi, 7, I-67100 L'Aquila, Italy", "Department of Physics, Kansas State University, 116 Cardwell Hall, Manhattan, Kansas 66506, USA", "Dipartimento di Fisica, Università di Napoli \"Federico II\", Via Cinthia, I-80126 Napoli, Italy; INFN Sezione di Napoli, Monte Sant'Angelo, Via Cinthia, I-80126 Napoli, Italy; Instituto de Ciencias Nucleares, Universidad Nacional Autónoma de México, México City, Districto Federal 04510, Mexico", "Department of Physics, Kansas State University, 116 Cardwell Hall, Manhattan, Kansas 66506, USA" ]	[ "2014AnP...526..309C", "2014AstRv...9d...6M", "2014arXiv1411.2350C", "2015CQGra..32m5007V", "2015IJMPD..2441002C", "2015JCAP...09..041M", "2015JCAP...09..045V", "2015PhLB..747....1S", "2015PhRvD..91l4037C", "2015PhRvD..91l4048P", "2016EPJP..131..256S", "2016Entrp..18...94K", "2016GReGr..48...53S", "2016IJMPD..2530010C", "2016JCAP...05..014M", "2016JCAP...07..032A", "2016JCAP...11..052M", "2016JCAP...12..042D", "2016PhRvD..94h3525D", "2017ApJ...835...26F", "2017ApJ...835...86C", "2017ApJ...843...65R", "2017ApJ...850..183Z", "2017EPJC...77..480M", "2017EPJP..132...64R", "2017MPLA...3250105S", "2017PhRvD..96j3534R", "2017arXiv170602537R", "2017arXiv171207569M", "2018ApJ...856....3Y", "2018ApJ...868...83P", "2018EL....12139001S", "2018EPJC...78..213L", "2018EPJC...78..495S", "2018EPJC...78..736S", "2018EPJC...78..862A", "2018GReGr..50...53C", "2018Galax...6...28D", "2018IJMPA..3350207S", "2018MNRAS.476.3924C", "2018MNRAS.480..759R", "2018MPLA...3350193S", "2018PhRvD..97f3503A", "2018PhRvD..97l3525D", "2018PhRvD..98b4009C", "2018PhRvD..98f3533M", "2018arXiv180502854M", "2018arXiv180904043H", "2019Ap&SS.364..118T", "2019Ap&SS.364..134P", "2019Ap&SS.364..135B", "2019CaJPh..97..966H", "2019EPJC...79..175P", "2019EPJC...79..267T", "2019IJMPD..2830016C", "2019JCAP...02..053A", "2019JPhCS1344a2004S", "2019LRR....22....1I", "2019NewA...7301284G", "2019Prama..93...89G", "2020Ap&SS.365..103B", "2020ApJ...898...82M", "2020EPJP..135....1C", "2020GReGr..52...13A", "2020GReGr..52...32M", "2020GrCo...26..144S", "2020IJGMM..1750194H", "2020IJMPD..2950097A", "2020NewA...7701355M", "2020PDU....3000722T", "2020PhyS...95k5001T", "2020Symm...12.1118M", "2020arXiv200604339B", "2021ApJ...908...84V", "2021EPJC...81..237O", "2021ForPh..6900007G", "2021IJGMM..1850095H", "2021IJMPA..3650148S", "2021IJMPD..3040005T", "2021MPLA...3630014S", "2021NewA...8401465H", "2021NewA...8601567A", "2021PhLB..82336786M", "2021PhRvD.103f4063K", "2021PhRvD.104f4044Y", "2021ResPh..2303863V", "2021arXiv210410354R", "2021arXiv211002071N", "2022ApJ...927..164B", "2022ChJPh..80..261B", "2022EL....13739001O", "2022EPJC...82..765R", "2022EPJP..137.1294V", "2022NuPhB.97815738K", "2022PDU....3601033N", "2022PDU....3801126K", "2022PDU....3801137H", "2022PhLB..83437420K", "2022PhRvD.105b4006K", "2022PhRvD.105d4038K", "2022Univ....8..484O", "2022arXiv220409496N", "2023A&C....4500768D", "2023AnP...53500161N", "2023AnP...53500178E", "2023Ap&SS.368...32V", "2023ChPhC..47k5107M", "2023EPJP..138..301S", "2023EPJP..138..469F", "2023GReGr..55...95S", "2023GrCo...29..468T", "2023InJPh..97.2217D", "2023JApA...44....6K", "2023MNRAS.523.4938M", "2023MPLA...3850093R", "2023PDU....3901164L", "2023PDU....4001174N", "2023RAA....23i5001D", "2023ResPh..5507166K", "2023arXiv230301985N", "2023arXiv230903065L", "2024A&C....4600789M", "2024Ap&SS.369...50B", "2024BlgAJ..40...95A", "2024ChJPh..87..665B", "2024EPJC...84..402D", "2024InJPh.tmp..228J", "2024PDU....4301402D", "2024arXiv240218083C", "2024arXiv240218997C", "2024arXiv240407070L", "2024arXiv240407313P", "2024arXiv240412707G" ]	[ "astronomy", "physics" ]	5	[ "04.50.-h", "04.20.Ex", "04.20.Cv", "98.80.Jk", "Higher-dimensional gravity and other theories of gravity", "Initial value problem existence and uniqueness of solutions", "Fundamental problems and general formalism", "Mathematical and relativistic aspects of cosmology", "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]	[ "1953JChPh..21.1087M", "1973JSP.....8....1M", "1984ApJ...284..439P", "1988ApJ...325L..17P", "1988PhRvD..37.3406R", "2000ApJ...532..109P", "2001ApJ...549....1G", "2001ApJ...553...39P", "2001ApJ...559....9P", "2003JETPL..77..201S", "2003PASP..115.1269C", "2003RvMP...75..559P", "2004ApJ...607L..71C", "2004CQGra..21.2603V", "2005ASPC..339...27B", "2005GReGr..37.1541V", "2005PhRvD..71d3503C", "2005PhRvD..71l3001S", "2005PhRvD..72d4022C", "2005PhRvD..72h3505O", "2006MNRAS.368..371W", "2006MPLA...21.1083S", "2007CQGra..24.5985C", "2007JETPL..86..157S", "2008MNRAS.390..210C", "2008PhRvD..77b3507T", "2008PhRvD..78f3501C", "2008PhRvD..78f3504C", "2008cosm.book.....W", "2009AdAst2009E...7C", "2009PhRvD..79d4001C", "2009PhRvD..79d7301C", "2009PhRvD..80f7301M", "2010ApJ...714.1347S", "2010JCAP...02..008S", "2010PhRvD..82j3526B", "2011MNRAS.410.1993S", "2011MPLA...26.1459L", "2011PASP..123.1127C", "2011PhR...505...59N", "2011PhR...509..167C", "2011PhRvD..83d3501A", "2011adco.book..355L", "2012Ap&SS.338..345L", "2012Ap&SS.342..155B", "2012ApJ...746...85S", "2012ApJ...753...97W", "2012ApJ...760...19P", "2012JCAP...08..006M", "2012MNRAS.423.2503S", "2012MNRAS.425..405B", "2012MNRAS.426..226C", "2012MNRAS.426.1396D", "2012PhRvD..85d3520X", "2012PhRvD..86d3520C", "2012PhRvD..86l3516A", "2012arXiv1205.0548B", "2013A&A...552A..96B", "2013A&ARv..21...67B", "2013ApJ...766L...7F", "2013ApJ...769..133N", "2013ApJS..208...19H", "2013CQGra..30t5003H", "2013EPJC...73.2497C", "2013IAUS..289..319C", "2013IJMPD..2250059L", "2013JCAP...11..020B", "2013MPLA...2850080L", "2013PhLB..718.1194A", "2013PhLB..726...72F", "2013PhLB..727....8Z", "2013PhRvD..87b3532A", "2013PhRvD..87d4012A", "2013PhRvD..87h3001A", "2013PhRvD..88d3522W", "2013PhRvD..88d4043T", "2013PhRvD..88f3531P", "2013PhRvD..88f3534Z", "2013PhRvD..88h3503F", "2013PhRvD..88j3004B", "2013PhRvD..88l3513P", "2013RAA....13.1013W", "2013arXiv1301.4044A", "2013arXiv1309.4133B", "2013arXiv1309.4188S", "2014A&A...571A..16P", "2014CQGra..31j5007G", "2014EPJP..129...22A", "2014FoPh...44..213W", "2014GReGr..46.1732P", "2014IJMPD..2350012L", "2014MNRAS.440.1138E", "2014PhLB..732..330C", "2014PhRvD..89h3518B", "2014PhRvD..89j3506G", "2014PhRvD..90b3006P", "2014RAA....14.1221Z", "2014arXiv1401.0046M", "2015Ap&SS.357...11F", "2015PNAS..11212249W" ]	[ "10.1103/PhysRevD.90.044016", "10.48550/arXiv.1403.1421" ]	1403	1403.1421_arXiv.txt	The inclusion of a cosmological constant $\Lambda$ in Einstein's equations is arguably the simplest way to produce accelerated cosmological expansion. The corresponding cosmological model, namely the $\Lambda$CDM model \cite{Peebles84}, predicts a currently accelerating cosmological expansion and is in fairly good agreement with current observations \cite{Sami}.\footnote{In this model, the current cosmological energy budget is dominated by $\Lambda$, with cold dark matter (CDM) in second place, and baryons a distant third. The $\Lambda$CDM model assumes the simplest form of CDM, which might be in conflict with some observations on structure formation \cite{lambdaobs}.} However, the $\Lambda$CDM model has some puzzling features \cite{Sola}. The first puzzle is that both matter and $\Lambda$ energy densities are comparable in order of magnitude today. The second puzzle is the huge difference between the observed $\Lambda$ and the corresponding quantum field theory naively-computed value. Perhaps these puzzles mean that the standard $\Lambda$CDM model is only a limiting case of more complete, and less puzzling, cosmological model. In such a model the role played by $\Lambda$ in the $\Lambda$CDM model might be generalized to another, more complex but still unknown, substance, often dubbed (dynamical) dark energy \cite{darkenergy}. The cosmological constant can be considered to be a fluid with equation of state $p_{\Lambda}=-\rho_{\Lambda}$ relating the time-independent energy density $\rho_{\Lambda}$ and pressure $p_{\Lambda}$. A fluid with equation of state $p_X=\omega_X\rho_X$ relating the $X$-fluid energy density $\rho_X$ and pressure $p_X$ is a very simple (but incomplete \cite{Podariu}) model of dynamical dark energy density. The $\phi$CDM model \cite{PeeblesandRatra}, where dynamical dark energy density is modeled by a scalar field $\phi$ with potential energy density $V(\phi)$ is a complete and consistent model. Many models have followed this one during the last quarter century, and an evolving dark energy fluid may be responsible for the late time acceleration, if the corresponding equation of state parameter, $\omega$, is within the interval $-1<\omega\simeq -0.75$, for redshift $z\ll1$ \cite{mico}. However, none of them has managed to satisfactorily clarify the physical origin of dark energy, thus a definite explanation of the accelerated cosmological expansion remains elusive. A major shortcoming of the dark energy paradigm is the lack of more fundamental first principles (less phenomenological), motivation for dark energy \cite{notes}. Another possibility, is to, instead, ascribe the observed accelerating cosmological expansion to modified gravity, a generalization of general relativity \cite{coppa4}. This is the underlying philosophy of modified theories of gravity, which extend general relativity by adding further curvature invariants to the Einstein-Hilbert action \cite{coppa45}. Much discussed extensions of standard cosmology arise when the Ricci scalar is replaced by an analytic function of curvature, $f(\mathcal{R})$ \cite{curvatura}. Here the action is \begin{equation}\label{azione} {\cal{A}} = \int{d^4x \sqrt{-g} \left [ f(\mathcal{R}) + {\cal{L}}_{m}\right]}\,, \end{equation} where ${\cal{L}}_{m}$ is the matter Lagrangian, involving generally both baryons and cold dark matter and $g$ is the determinant of the metric tensor. By varying the action with respect to the metric tensor $g_{\mu\nu}$, one infers the field equations \begin{eqnarray}\label{filedeqs} \mathcal{R}_{\mu \nu}f^{'}(\mathcal{R})-\frac{1}{2}f(\mathcal{R})g_{\mu \nu}-(\nabla_{\mu}\nabla_{\nu}-g_{\mu \nu}\nabla_{\alpha}\nabla^{\alpha})f^{'}(\mathcal{R})\nonumber \\ =8\pi T_{\mu \nu}\,. \end{eqnarray} Here a prime denotes a derivative with respect to $\mathcal{R}$, $\mathcal{R}_{\mu \nu}$ is the Ricci tensor, $T_{\mu\nu}$ is the standard energy-momentum tensor and we have set the Newtonian gravitational constant and the speed of light to unity $G=c=1$. Dark energy can therefore be considered to be a geometrical fluid that adds to the conventional stress-energy tensor, if one want to interpret this equation in the context of general relativity \cite{geometric}. In doing so, the task of determining the dark energy equation of state is replaced by trying to understand which $f(\mathcal{R})$ better fits current data. It has been argued that viable candidates of $f(\mathcal{R})$ are those that reduce to the $\Lambda$CDM model at $z\ll1$ \cite{beng}. This guarantees fairly good agreement with present observations, permitting us to ease the experimental problems associated with wrong choices of $f(\mathcal{R})$. In the $\Lambda$CDM model, and in dynamical dark energy models, the dark energy density has only recently come to dominate the cosmological energy budget and so accelerate the cosmological expansion. Earlier on dark matter dominated, resulting in decelerating cosmological expansion. Recent cosmological measurements have lead to the first believable estimate of the decelerating-accelerating transition redshift $z_{\mathrm da}$, \cite{ratra13}, of order 0.75. Thus there is now some observational support for the dark energy idea at redshifts approaching unity. Therefore it is of interest to see whether such data are also consistent with $f(\mathcal{R})$ gravity models. In this work we determine observationally viable $f(\mathcal{R})$ models by assuming the cosmological principle and using cosmography to constrain parameters \cite{cosmography1}. From the Taylor series expansion of the scale factor $a(t)$, cosmography can be used to numerically bound late time measurable quantities, e.g., the acceleration parameter, the jerk parameter, the snap, and so forth, \cite{cosmography2}. Possible departures from the standard $\Lambda$CDM model could be determined through the use of cosmography, which represents a tool to pick the most viable class of $f(\mathcal{R})$ models \cite{cosmography3}. To this end, cosmography allows us to relate the expanded quantities of interest, i.e. the Hubble rate, luminosity distances, magnitudes, and so forth, in terms of observables \cite{cosmography4}. In this paper, we show that a particular Hubble rate, derived from a viable class of $f(\mathcal R)$ models, predicts a transition from decelerated to accelerated cosmological expansion at a transition redshift which is in fairly good agreement with the cosmographic series. This is an extension of the standard $\Lambda$CDM model with a logarithmic term that mimics the effect of the dark energy as a smoothly varying function of the redshift $z$. The corresponding acceleration parameter changes sign around $z\sim 0.75$, in agreement with the recent measurement \cite{ratra13}. We also rewrite $f(\mathcal{R})$ as a function of $z$, as a series in the scale factor, i.e. $a\equiv(1+z)^{-1}$. This allows us to describe the curvature dark energy fluid in terms of the more practical redshift variable. In turn, we determine observational bounds on the cosmographic series and on the expanded $f(\mathcal R)$, by combining the most recent Union 2.1 supernova apparent magnitude compilation \cite{suzuky} and Hubble rate measurements in the interval $z\in[0,2.8]$ \cite{ratra13,ultima1}, through the use of Monte Carlo analyses using the the Metropolis algorithm \cite{metro}. We obtain our fits by using ROOT \cite{root} and BAT \cite{bat}. Our paper is structured as follows: In Sec.\ II we set the initial conditions on cosmological observables, through the use of cosmography. These initial conditions are useful when determining constraints on $f(z)$ at low redshift. In Sec.\ III we relate these initial conditions to the modified Friedmann equations. In Sec.\ IV we describe the corresponding cosmological model, inferred from numerically solving the modified Friedmann equations, with the numerical bounds inferred from cosmography. Furthermore, we describe the obtained transition redshift and we discuss the numerical results by comparing our model to $\Lambda$CDM. In Sec.\ V model parameters are computed by using supernovae apparent magnitude and Hubble parameter measurements. In Sec.\ VI we provide a physical interpretation of the free parameters of the model, relating them to derivatives of the Ricci scalar at the present time. Finally, Sec.\ VII is devoted to our conclusions.	We have numerically analyzed a class of $f\left(\mathcal{R}\right)$ gravity models which reduce to $\Lambda$CDM at $z\simeq0$. Deviations emerge at third order in the $f(\mathcal R)$ Taylor expansion. As present epoch constraints we adopt the cosmographic series, i.e. the series of measurable coefficients derived by expanding the luminosity distance and comparing it with data. We therefore inferred cosmographic bounds on the test function $f(z)$ which reproduces the observed low redshift cosmological behavior. Since cosmography allows for a determination of model independent constraints on $f(z)$ and derivatives, we used a Taylor expansion of $f(z)$ in terms of $a(t)={1}/{(1+z)}$, which fairly well approximates the Friedmann equations in the range $z\leq 2$. We found good agreement, with small departures at $z\ll1$ and $z\sim2$, for the range of parameters $\tilde f_0\sim-10$, $\tilde f_1\sim 7$, $\tilde f_2\sim-3.7$, $\sigma_1=1$ and $\sigma_2=2$, which are compatible with the initial conditions defined by cosmography. Such departures lead to possible logarithmic corrections of the conventional Hubble rate, showing an evolving dark energy term different from the cosmological constant. We demonstrated that this model has a transition redshift in a range compatible with measurements \cite{ratra13}. To this end, cosmological constraints on the model were determined using a Monte Carlo approach based on the Metropolis algorithm. Our model passes all the cosmological tests, showing that the obtained curvature dark energy is compatible with observations. We implemented different priors on the fitting parameters, and in particular, we fixed $H_0$ to the \textit{Planck} value first and then to a numerical value obtained by fitting the first-order luminosity distance $d_\mathcal L$ to the supernova data in the interval $z\in[0,0.36]$. In general, results seem to indicate slightly less negative acceleration parameters with non-conclusive results on the variation of acceleration, namely the jerk parameter. Using these results we provided a self consistent explanation of the free parameters of the model, showing that they could be related to the terms of the Taylor series of $f(\mathcal{R})$. In doing so, by comparing our results with PPN approximations, we found that $\alpha$ and $\beta$ could be related to third-order PPN parameters. Future investigations will be devoted to better constraining the logarithmic correction due to $f(\mathcal R)$.	14	3	1403.1421	We examine the observational viability of a class of f(R) gravity cosmological models. Particular attention is devoted to constraints from the recent observational determination of the redshift of the cosmological deceleration-acceleration transition. Making use of the fact that the Ricci scalar is a function of redshift z in these models, R =R(z), and so is f(z), we use cosmography to relate a f(z) test function evaluated at higher z to late-time cosmographic bounds. First, we consider a model-independent procedure to build up a numerical f(z) by requiring that at z=0 the corresponding cosmological model reduces to standard ΛCDM. We then infer late-time observational constraints on f(z) in terms of bounds on the Taylor expansion cosmographic coefficients. In doing so we parametrize possible departures from the standard ΛCDM model in terms of a two-parameter logarithmic correction. The physical meaning of the two parameters is also discussed in terms of the post-Newtonian approximation. Second, we provide numerical estimates of the cosmographic series terms by using type Ia supernova apparent magnitude data and Hubble parameter measurements. Finally, we use these estimates to bound the two parameters of the logarithmic correction. We find that the deceleration parameter in our model changes sign at a redshift consistent with what is observed.	false	[ "gravity cosmological models", "Hubble parameter measurements", "redshift z", "type Ia supernova apparent magnitude data", "terms", "f(z", "standard ΛCDM", "higher z", "post-Newtonian", "R =", "the corresponding cosmological model", "late-time cosmographic bounds", "the Taylor expansion cosmographic coefficients", "the cosmographic series terms", "the standard ΛCDM model", "bounds", "numerical estimates", "the post-Newtonian approximation", "late-time observational constraints", "Hubble" ]	10.435131	1.147866	89
745737	[ "Finn, Charles W.", "Morris, Simon L.", "Crighton, Neil H. M.", "Hamann, Fred", "Done, Chris", "Theuns, Tom", "Fumagalli, Michele", "Tejos, Nicolas", "Worseck, Gabor" ]	2014MNRAS.440.3317F	[ "A compact, metal-rich, kpc-scale outflow in FBQS J0209-0438: detailed diagnostics from HST/COS extreme UV observations" ]	28	[ "Department of Physics, Durham University, South Road, Durham DH1 3LE, UK; Institute for Computational Cosmology, Department of Physics, Durham University, South Road, Durham DH1 3LE, UK", "Department of Physics, Durham University, South Road, Durham DH1 3LE, UK", "Max Planck Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany; Centre for Astrophysics and Supercomputing, Swinburne University of Technology, PO Box 218, Victoria 3122, Australia", "Department of Astronomy, University of Florida, 211 Bryant Space Sciences Center, Gainesville, FL 32611-2055, USA", "Department of Physics, Durham University, South Road, Durham DH1 3LE, UK", "Institute for Computational Cosmology, Department of Physics, Durham University, South Road, Durham DH1 3LE, UK; Department of Physics, University of Antwerp, Campus Groenenborger, Groenenborgerlaan 171, B-2020 Antwerp, Belgium", "Carnegie Observatories, 813 Santa Barbara Street, Pasadena, CA 91101, USA; Department of Astrophysics, Princeton University, Princeton, NJ 08544-1001, USA", "Department of Physics, Durham University, South Road, Durham DH1 3LE, UK; Department of Astronomy and Astrophysics, UCO/Lick Observatory, University of California, 1156 High Street, Santa Cruz, CA 95064, USA", "Max Planck Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany" ]	[ "2014ApJ...794...75S", "2015MNRAS.446...18C", "2015MNRAS.449.2174C", "2016ApJ...830..104P", "2016ApJ...832..189Q", "2016ApJ...833L..26G", "2016MNRAS.455.2662T", "2016MNRAS.460..590F", "2017MNRAS.465..358C", "2018ApJ...857...60A", "2018ApJ...858...39X", "2018ApJ...865...90M", "2018ApJ...866....7L", "2018MNRAS.473.5407M", "2018MNRAS.481.3865C", "2019ApJ...876..105X", "2019ApJ...877...72T", "2019ApJ...879...27L", "2019ApJ...887...78C", "2019MNRAS.485.3409G", "2019MNRAS.487.5041H", "2020ApJS..247...37A", "2020ApJS..247...38X", "2020ApJS..247...41M", "2021ApJ...914...13T", "2022ApJ...937...74C", "2022ApJ...940L..40J", "2022Atoms..10..131R" ]	[ "astronomy" ]	16	[ "galaxies: active", "quasars: absorption lines", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1968ApJ...152..701B", "1975ApJ...198...95G", "1975ApJ...202..296W", "1979ApJ...229..891Y", "1979ApJ...234...33W", "1980ApJ...236..577B", "1982ApJ...252...54S", "1982MNRAS.198...91C", "1986ApJ...303..336G", "1986ApJ...307..504F", "1986ApJ...310...40M", "1989ApJ...345..245C", "1990ARA&A..28..215D", "1991ApJ...373...23W", "1991ApJ...379...52W", "1993ApJ...418...11H", "1994A&A...291...29P", "1994ApJ...434..493M", "1995ApJ...443..606H", "1995ApJ...447..512K", "1995ApJ...455..448D", "1996ApJS..102..239T", "1997ADNDT..65....1V", "1997AJ....113..136B", "1997ApJS..109..279H", "1998A&A...331L...1S", "1998A&AS..133..403M", "1998ApJ...503L..23M", "1998ApJ...509..132T", "1999A&A...345...73P", "1999A&A...349..389V", "1999AJ....117.2594G", "1999ARA&A..37..487H", "1999ApJ...512..184N", "1999ApJ...516...27A", "1999ApJ...518..672S", "2000A&A...357..414S", "2000ApJ...528..617S", "2000ApJ...528..637B", "2000ApJ...536..101H", "2000ApJ...542..914W", "2001A&A...373..816S", "2001ApJ...550..142H", "2001ApJ...559..507S", "2002ApJ...574..643K", "2002ApJ...576....1A", "2002ApJ...579..530W", "2002ApJ...580...54D", "2003ARA&A..41..117C", "2003ApJ...596L..27K", "2003ApJ...599...38B", "2004A&A...417..515V", "2004ASPC..311..203H", "2004ApJ...608...62S", "2004MNRAS.351..976D", "2004PASP..116..362C", "2005A&A...440..775K", "2005ARA&A..43..769V", "2005ApJ...623...85G", "2005Natur.433..604D", "2006AJ....131.1163S", "2006ApJ...639..590F", "2006ApJ...646..742G", "2006ApJS..165....1T", "2006ApJS..166..634T", "2006ApJS..167..334B", "2006MNRAS.366..449F", "2006MNRAS.370..645B", "2007ApJ...658..829A", "2007ApJS..171....1M", "2008ApJ...681..954A", "2008MNRAS.386..935F", "2008MNRAS.386.2055N", "2008MNRAS.388..227W", "2009ApJ...706..525M", "2010A&A...516A..89R", "2010ApJ...722..642O", "2010MNRAS.401....7H", "2010MNRAS.404.1464M", "2011Ap&SS.335..257O", "2011ApJ...737...26N", "2011ApJ...739....7E", "2011MNRAS.411.2653H", "2011MNRAS.418..357B", "2012A&A...544A..33A", "2012ApJ...744...60G", "2012ApJ...751..106J", "2012ApJ...751..107B", "2012ApJ...752..162S", "2012MNRAS.420.1347F", "2012MNRAS.420.1825J", "2012MNRAS.420.1848D", "2012MNRAS.423.2879V", "2012MNRAS.424L..59M", "2012PASP..124..830K", "2013A&A...556A..94D", "2013A&A...558A..33A", "2013ApJ...764L..21F", "2013ApJ...771...64M", "2013ApJ...775...78F", "2013ApJS..208...17A", "2013MNRAS.429.1922C", "2013MNRAS.430.2650L", "2013MNRAS.431.2885M", "2013MNRAS.431L..93S", "2013MNRAS.434..748A", "2013MNRAS.434.1043O", "2013MNRAS.434.1063O", "2013MNRAS.435..133H", "2013MNRAS.435.1233G", "2013MNRAS.436.3286A", "2013PASP..125..270S", "2013RMxAA..49..137F", "2014ApJS..211...17A", "2014MNRAS.437.2017T" ]	[ "10.1093/mnras/stu518", "10.48550/arXiv.1403.3338" ]	1403	1403.3338_arXiv.txt	\label{introduction} \ac{aals} seen in \ac{qso} spectra offer a unique physical perspective on the gaseous environments in the vicinity of \ac{qso}s. Many \ac{aals} are thought to arise in material that has been ejected from a region close to the \ac{smbh} \cite[within a few pc, e.g.][]{Nestor:2008gb,Wild:2008hn,Muzahid:2012kl}. The resulting outflows might play a major role in the quenching of star formation and in regulating the growth of \ac{smbh}s \citep{Silk:1998up,King:2003gt,Scannapieco:2004es,DiMatteo:2005hl,Ostriker:2010ik,Hopkins:2010cf}. Some may arise from material ejected in supernova explosions \cite[e.g.][]{Veilleux:2005ec}. In addition, some \ac{aals} appear to probe gas that is part of the host galaxy halo \cite[e.g.][]{Williams:1975ge,Sargent:1982kl,Morris:1986fa,Tripp:1996ge,Hamann:2001eg,DOdorico:2004jd}. In some cases, this gas may eventually condense in the disc to form new generations of stars via Galactic fountain processes \citep{Bregman:1980gn,Fraternali:2013dd}. The balance of gas accretion and outflow shapes the galaxy luminosity function and drives the evolution of galaxies \cite[e.g.][]{Benson:2003ch,Bower:2006fj}. Observations of \ac{aals} therefore provide a detailed snapshot of these forces at work. Constraints on the metallicity of these absorbers also provides a direct measure of the star formation and chemical enrichment histories in the centres of active galaxies \citep{Hamann:1993jb,Hamann:1999ky}. \ac{aals} are loosely defined as having absorption redshifts within a few thousand \kms\ of the \ac{qso} emission redshift, and velocity widths of less than a few hundred \kms. These are narrow when compared to the so-called \ac{bals}, which have velocity widths and displacements from the \ac{qso} redshift that often exceed $10^4$~\kms \citep{Weymann:1979bt,Foltz:1986fg,Weymann:1991cn,Trump:2006ht}. The origin of the \ac{bals} is presumably in a wind, driven by accretion processes close to a \ac{smbh}. However, the exact origin of the \ac{aals} is far less clear. In addition, not all \ac{aals} are necessarily intrinsic to the \ac{qso} \cite[e.g.][]{Tripp:1996ge,Ganguly:2013hd}. Those that may be intrinsic show: (i) Absorption strength that is seen to vary on time-scales of around a year \cite[e.g.][]{Hamann:1995ff,Srianand:2001hh,Hall:2011ej,Vivek:2012hh}. (ii) Metallicities $\gtrsim Z_{\odot}$ \cite[e.g.][]{Petitjean:1994ti,Hamann:1997iu,Muzahid:2013dm}. (iii) Partial coverage of the \ac{qso} accretion disc continuum and/or \ac{blr} \cite[e.g.][]{Barlow:1997eb,Srianand:1999hj,Gabel:2006fw,Arav:2008fa}. (iv) The presence of excited fine structure lines \cite[e.g.][]{Morris:1986fa,Srianand:2000tq,Hamann:2001eg,Edmonds:2011fz}. These properties are rarely seen in intervening absorption-line systems \cite[see][for an exceptional case]{Balashev:2011cc}, and so \ac{aals} with these properties are believed to trace gas that originates near the \ac{qso}, or in the halo of the host galaxy. \ac{aals} have been observed in optical, \ac{uv} and X-ray spectra of local \ac{agn} and \ac{qso}s, with the X-ray observations often revealing a plethora of absorption lines and K-shell absorption edges from species with ionization potentials of a few hundred eV (e.g. \ion{O}{7} and \ion{O}{8}). Collectively, these are usually referred to as `warm absorbers' \cite[in $\sim 50$\% of Seyfert galaxies;][]{Crenshaw:2003hz}. Many authors have suggested that the presence of warm absorbers is correlated with the detection of \ac{aals} and \ac{bals} in optical and \ac{uv} spectra, usually through species like \ion{C}{3}, \ion{C}{4} and \ion{N}{5}, with ionization potentials $\lesssim 100~\textrm{eV}$ \cite[e.g.][]{Mathur:1994ds,Mathur:1998fo,Brandt:2000de,Kaspi:2002cr,Arav:2007dx,DiGesu:2013cf}. However, at present, it is not clear whether these correlations imply a physical connection between the gas clouds traced by these ions \cite[see, for example,][]{Srianand:2000jd,Hamann:2000bi,Hamann:2013cq}. To better understand the nature of associated absorbing clouds, more observations of the most highly ionized \ac{uv} species (ionization potentials $> 100~\textrm{eV}$) are required, so that the ionization structure of the absorbing gas can be more extensively characterised. At low redshift, observations of many \ac{uv} ions are impossible due to the presence of Galactic Lyman limit absorption (the relevant transitions have rest-frame wavelengths $< 912$~\AA). Observations in the optical, which are limited to high redshifts, are complicated by contamination from the Lyman alpha forest, together with a higher incidence rate of Lyman limit systems \citep{Fumagalli:2013ee}. At intermediate redshifts $0.5 \lesssim z \lesssim 1.5$, the problem of Galactic absorption is virtually eliminated, and the Lyman alpha forest contamination is less severe, making this a profitable redshift range to study highly ionized \ac{aals}. Observations must be conducted in the \ac{fuv}, and with the advent of the \ac{cos} on-board the \ac{hst}, hundreds of \ac{qso}s are now observable in this wavelength regime, thanks largely to a sensitivity more than ten times that of the previous generation medium resolution \ac{uv} spectrograph \citep{Green:2012dj}. Together with the \ac{nuv} modes of \ac{cos}, \ac{aals} with ionization parameters of a few, to a few hundred eV are accessible. Detailed diagnostics on the ionization structure of associated gas clouds are thus available in a large number of \ac{qso}s for the first time. In addition, coverage of strong transitions due to fine-structure excited states in ions such as \ion{O}{4} and \ion{O}{5} (see fig. 1 in \citet{Arav:2013be} for a full summary) provide powerful density diagnostics in highly ionized associated gas clouds, which provide crucial constraints on the physical conditions in and around the absorbing regions. In this paper, we present observations of the radio-loud \ac{qso} FBQS J0209-0438 obtained with \ac{cos}. This \ac{qso} was targeted as part of a larger programme of observations to study two-point correlation statistics between \ac{igm} absorbers and galaxies at $z \lesssim 1$ \cite[PID 12264, PI: S.L. Morris;][]{2014MNRAS.437.2017T}. A highly ionized system of \ac{aals} is present, with complex velocity structure and an overall velocity width $\approx 600$~\kms. We also report the detection of absorption due to the fine-structure \ion{O}{4}* transition. A summary of the observations and data reduction, together with a characterisation of the rest-frame \ac{qso} \ac{sed} is presented in \Cref{observations}. A complete analysis of the \ac{aals}; their covering fractions, column densities and line widths is presented in \Cref{analysis}. In \Cref{properties} we present the results of photoionization and collisional ionization models in an attempt to characterise the physical properties of the gas. In particular we examine the ionization state, metallicity and density of the gas, and use these properties to put constraints on the absorbing geometry and distance from the \ac{qso}. In \Cref{discussion} we present a discussion of these results and draw conclusions.	\label{discussion} The main results from the data analysis and photoionization/collisional ionization equilibrium models are as follows: \begin{enumerate} \item{The gas in the least ionized AAL region is predominantly photoionized.} \item{Under photoionization equilibrium, multiple ionization parameters are required to reproduce the column density ratios seen in the data.} \item{Based on the observed column densities and ionization fractions implied from characteristic ionization parameters, incorporating uncertainties in the shape of the \ac{qso} \ac{sed}, the gas metallicity is conservatively $0 \lesssim [\textrm{O} / \textrm{H}] \lesssim 1$.} \item{Given the range in possible gas metallicity, the total hydrogen column density in each velocity component is $10^{17} \lesssim N_{\textrm{H}} \lesssim 10^{18.5}~\textrm{cm}^{-2}$ in the least ionized gas (with slightly smaller values for velocity component $v_8$) and $10^{18.5} \lesssim N_{\textrm{H}} \lesssim 10^{20}~\textrm{cm}^{-2}$ through the most highly ionized gas we detect.} \item{Taking the column density ratio between \ion{O}{4}* and \ion{O}{4}, assuming the fine structure excited states are populated mostly due to collisions with electrons, the total hydrogen density in the gas traced by these ions is found to be $\log (n_{\textrm{H}} / \textrm{cm}^{-3}) \sim 3$ for solar metallicity.} \item{Given the total hydrogen density, and the plausible range in ionization parameter for the gas traced by \ion{O}{4} and \ion{O}{4}, the distance to the absorbing clouds from the centre of the \ac{qso} is found to be $2.3 \lesssim R \lesssim 6.0~\textrm{kpc}$. An empirically identified, shared velocity structure amongst all ions, suggests this distance determination is likely to be representative for the AAL region as a whole.} \item{Under photoionization equilibrium, the total hydrogen density in the most highly ionized AAL gas is found to be two orders of magnitude lower than that implied for the low ionization gas. Alternatively, models in which this gas is under \ac{cie} allow for densities to be similar across the AAL region probed by these data.} \item{The ratio $N_{\textrm{H}} / n_{\textrm{H}}$ sets limits on the absorption pathlength through the least and mostly highly ionized regions of $10^{-4.5} \lesssim l_{\textrm{abs}} \lesssim 10^{-3}~\textrm{pc}$ and $0.1 \lesssim l_{\textrm{abs}} \lesssim 1~\textrm{pc}$ respectively in each velocity component.} \item{Covering fractions less than unity (in all cases where they can be reliably measured), suggest that the continuum region is only partially covered, requiring clouds with transverse sizes $l_{\textrm{trans}} \lesssim 10^{-2.5}~\textrm{pc}$.} \end{enumerate} In summary, the analysis of the previous sections has revealed the presence of metal enriched (to at least solar), highly ionized gas clouds a few kpc from the centre of Q0209 that are likely to be very small (sub-pc scale). In the following sections we place these results in a wider context, and speculate on the origins and fate of the absorbing gas. For simplicity, we shall speak of two, co-spatial regions in ionization equilibrium: a low ionization, photoionized region with $\log U \lesssim -1$, and a high ionization region with $\log U \gtrsim -1$ if photoionized, or temperatures $T \gtrsim 10^{5.5}~\textrm{K}$ if collisionally ionized. \subsection{Gas structure and dynamics} \label{structure_dynamics} A redshift measurement for Q0209 of $z_{\textrm{QSO}} = 1.13194 \pm 0.00001$ implies that the AAL gas is mostly outflowing from the \ac{qso} with velocities up to $\sim 400$~\kms\ (see \Cref{aals}). This is unusually small, compared to the majority of the \ac{aals} and \ac{bals} in the literature with high ionization species such as \ion{Ne}{8} and \ion{Mg}{10}, which are typically outflowing with velocities closer to a few thousand or few tens of thousand \kms\ \cite[e.g.][]{Hamann:1995ff,Telfer:1998gx,Petitjean:1999vh,Arav:1999ka,Muzahid:2012kl,Muzahid:2013dm}, although see \cite{Hamann:2000bi} for a more similar example. If we assume that the gas is moving with a constant radial velocity $v$ and originates close to the \ac{smbh}, then the time-scale for reaching its current radius $R$ is at least \begin{equation} t \approx 10^7 \left(\frac{R}{2.3~\textrm{kpc}}\right)\left(\frac{200~\textrm{km~s}^{-1}}{v}\right)~\textrm{yrs}. \label{flow_time} \end{equation} Different velocity components in the AAL gas are moving at different speeds, so the overall region should possess an appreciable radial thickness after a time $t$, even though we derive densities and ionization parameters that are consistent with one another across the different velocity components. Given the small cloud sizes in the low ionization gas, a key question is how long they are expected to survive. The free-fall time-scale for these clouds can approximately be expressed as \begin{equation} t_{\textrm{ff}} \equiv \frac{1}{\sqrt{G \rho}} \sim 1.0 \times 10^{15}~\textrm{s} \left(\frac{n_{\textrm{H}}}{\textrm{cm}^{-3}}\right)^{-1/2}, \end{equation} where $G$ is the gravitational constant, $\rho$ is the gas density, assuming that all of the mass is baryonic, and setting the mass fraction in helium to 0.28 (assuming solar abundances). In addition, for a characteristic cloud size $l$, the sound crossing time in a highly ionized plasma can be approximated by \begin{equation} t_{\textrm{sc}} \equiv \frac{l}{c_{\textrm{s}}} \sim 2.1 \times 10^{15}~\textrm{s} \left(\frac{l}{\textrm{kpc}}\right) T_4^{-1/2}, \label{sc} \end{equation} where $c_{\textrm{s}}$ is the sound speed in an ideal monatomic gas and $T = T_4 \times 10^4~\textrm{K}$ \cite[e.g.][]{Schaye:2001dv}. For a stable cloud in hydrostatic equilibrium, $t_{\textrm{sc}} \sim t_{\textrm{ff}}$. We take a value of $n_{\textrm{H}} = 10^3~\textrm{cm}^{-2}$, and a value of $l = 10^{-6}~\textrm{kpc}$ (assuming $l_\textrm{abs} \approx l$). The photoionization models indicate that $T_4 \approx 2$ in this gas and so we find $t_{\textrm{ff}} \sim 3.2 \times 10^{13}~\textrm{s} \gg t_{\textrm{sc}} \sim 1.5 \times 10^9~\textrm{s}$. This implies that the clouds will expand on the sound crossing time-scale, so they should have lifetimes of $\lesssim 100$~years. This is considerably less than the characteristic flow time in \cref{flow_time}, and so the probability of observing these clouds at their implied distance from the \ac{qso} is extremely small in this case. The analysis presented above poses a problem, which may be overcome if the clouds are being held in pressure equilibrium. This may be a thermal pressure equilibrium with higher temperature, lower density, more highly ionized gas, equivalent to the statement $n_{\textrm{H}1} T_1 = n_{\textrm{H}2} T_2$, where $n_{\textrm{H}1}$, $T_1$ and $n_{\textrm{H}2}$, $T_2$ are the total hydrogen densities and temperatures of the low and high ionization regions respectively. From \Cref{properties}, under photoionization equilibrium we found that $n_{\textrm{H}2} \sim 10~\textrm{cm}^{-2}$, and these models also indicate that $T_2 \approx 6 \times 10^4~\textrm{K}$. In this case, $n_{\textrm{H}1} T_1 > n_{\textrm{H}2} T_2$, and the high ionization gas cannot pressure support the low ionization gas. If the former is collisionally ionized, we now have temperatures that differ by more than an order of magnitude. Densities in the high ionization region may be low enough to allow for pressure support. Nevertheless, the high ionization gas itself, accounting for the possibility that it is photoionized, should have a lifetime $\lesssim 10^5$ years, which is still short enough to suggest that this gas may also require pressure support from even more highly ionized gas that we do not detect in the UV, and which would require larger total column density, higher temperature and lower density. Massive galaxies are expected to host hot gas coronae, well within the implied location of the AAL region, with $T \sim 10^6~\textrm{K}$ and $n_{\textrm{H}} \sim 10^{-2}~\textrm{cm}^{-2}$ \cite[e.g.][]{White:1991in,Fukugita:2006dg}. Pressure from this external medium, together with additional pressure support from magnetic confinement \citep{deKool:1995gu} may help to alleviate the problems outlined above, although pressure supporting gas with varying internal pressure is clearly a complex issue. We note that the analysis above does not incorporate the effects of turbulence, which is almost certainly present given the observed line widths (see \Cref{column_densities}). In addition, gas outflowing from a \ac{qso} will likely encounter the \ac{ism} of the host galaxy on its journey out into the halo. At supersonic velocities, shocks will likely occur at the interface between the outflowing gas and the \ac{ism}, heating the gas close to this interface. The resulting mix of hot and cool gas creates instabilities that can destroy the clouds before they reach the halo \cite[see for example arguments in][]{FaucherGiguere:2012jk}. These authors suggest an alternative scenario, in which small clouds may be formed in situ from moderately dense \ac{ism} clouds within hot, recently shocked gas. These clouds become shredded by a passing blast wave and gain momentum from an accompanying shock. The resulting `cloudlets' in this model have sizes and densities comparable to those derived here, and can possess a range of velocities that may explain the multi-component velocity structure in the absorption profiles. Models such as these may offer a more promising route to explain the structure and dynamics of \ac{aals} with properties (density, cloud size, velocity structure, distance from the \ac{qso}) similar to those found in Q0209 \cite[e.g.][]{Petitjean:1999vh,Hamann:2000bi,Edmonds:2011fz,Borguet:2012ca,Arav:2013be,Muzahid:2013dm}. Any viable model must additionally reproduce the covering fractions seen in the present data. Covering fractions less than unity, and with little variation, are seen in ions spanning a range in ionization potential from a few tens to a few hundreds of eV, tracing gas with more than one possible ionization mechanism. Vastly differing absorption path lengths through the AAL region, as hinted at in the analysis of the previous sections, make it very difficult to account for the near constancy in covering fraction across all ions using simple geometrical models. We also note that these results differ from e.g. \citet{Hamann:2000bi} and \citet{Borguet:2012ca}, who find more complete coverage in high-ionization \ac{uv} transitions compared to those at lower ionization potentials. Covering fractions less than unity across our sample also go against general trends for more complete coverage with lower outflow velocities, as identified in the \ac{cos} sample presented by \citet{Muzahid:2013dm}. It is therefore clear that simple trends such as these may not produce robust predictions for individual systems, which further highlights the apparent complexity in these absorbers. \subsection{Are the AAL clouds out of equilibrium?} \label{equilibrium} Up to this point, our analysis and discussion has assumed that the AAL clouds are in ionization equilibrium. However, in general, absorbers may be out of equilibrium when close to an AGN, due to recombination time-scales that can be long compared to typical AGN lifetimes and duty cycles \cite[e.g.][]{1995ApJ...447..512K,Nicastro:1999fj,Arav:2012gv,Oppenheimer:2013cr,Oppenheimer:2013dx}. The resulting recombination lag can lead to situations where high ionization stages like \ion{O}{6}, \ion{Ne}{8} and \ion{Mg}{10} are enhanced relative to the expectation from equilibrium models. We examine these issues here. We define the the photoionization rate $\Gamma_{M_i}$ (s$^{-1}$) for a given ion $M_i$ as \begin{equation} \Gamma_{M_i} \equiv \int_{\nu_{0, M_i}}^{\infty} \frac{4 \pi J_{\nu}}{h \nu} \sigma_{M_i} (\nu)~\textrm{d}\nu. \label{photoionization_rate} \end{equation} Here $\nu$ is the frequency, $\nu_{0, M_i}$ is the ionization frequency, $J_{\nu}$ is the intensity of the \ac{qso} radiation field (erg s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$), $\sigma_{M_i}$ is the photoionization cross-section, and $h$ is Planck's constant. The recombination rate (s$^{-1}$) into an ion $M_i$ is given by \begin{equation} R_{M_i} \equiv \alpha_{M_i} n_e, \end{equation} where $n_e$ is the electron number density, and $\alpha_{M_i}$ is the temperature dependent recombination rate coefficient (cm$^3$ s$^{-1}$) for that ion. Finally, the collisional ionization rate for an ion $M_i$ is \begin{equation} C_{M_i} \equiv \beta_{M_i} n_e, \end{equation} where $\beta_{M_i}$ is the temperature dependent collisional ionization rate coefficient (cm$^3$~s$^{-1}$). Neglecting Auger ionization and charge transfer, the population in an ion $M_i$ is then \begin{align} \frac{\textrm{d}n_{M_i}}{\textrm{d}t} &= -n_{M_i} (\Gamma_{M_i} + R_{M_{i - 1}} + C_{M_i}) + n_{M_{i + 1}} R_{M_i} \nonumber \\ &\qquad {} + n_{M_{i - 1}} (\Gamma_{M_{i - 1}} + C_{M_{i - 1}}). \label{time_dependent_ionization} \end{align} Now suppose that an absorber is in photoionization equilibrium, i.e. $\textrm{d}n_{M_i} / \textrm{d}t = 0$, at time $t = 0$, but there is a sudden change in the ionizing flux, such that $\Gamma_{M_i}(t > 0) = (1 + \delta) \Gamma_{M_i}(t = 0)$, where $-1 \leq \delta \leq \infty$. Taking the collisional ionisation rate to be negligible (a reasonable approximation for a photoionized plasma), it can then be shown that the $e$-folding time-scale for change in the ionic fraction is given by \begin{equation} t_{\textrm{change}} = \left[ -\delta \alpha_{M_i} n_e \left( \frac{n_{M_{i + 1}}}{n_{M_i}} - \frac{\alpha_{M_{i - 1}}}{\alpha_{M_i}} \right) \right]^{-1}, \label{time-scale} \end{equation} \citep{Arav:2012gv}, where negative time-scales indicate a decrease in the ionic fraction, and positive time-scales indicate an increase. For changes in the ionizing flux within an order of magnitude ($0.1 < 1 + \delta < 10$), these time-scales are typically $\sim 10$ years for the densities implied by the \ion{O}{4} analysis of \Cref{density}, assuming $T \sim 10^4$~K. Since the gas densities in the AAL region are much higher than those typical of the diffuse \ac{igm} and \ac{cgm}, these time-scales are much shorter than the typical AGN lifetime \cite[$\sim~\textrm{Myr}$ time-scales; e.g.][]{Novak:2011gc}. Therefore, photoionization equilibrium might be a valid assumption in our case, so long as these time-scales are also short compared to the time-scale over which the \ac{qso} luminosity changes. However, we must also consider the dynamical evolution of the AAL clouds. In the previous section, we found that, unless the clouds traced by \ion{O}{4} are pressure supported by an external medium, they will expand on time-scales $\lesssim 100$ years. If this process is occurring, then the clouds may be entering a non-equilibrium state due to a recombination lag. To determine whether or not this scenario is likely, we numerically solve the coupled, time-dependent ionization equations (\cref{time_dependent_ionization}) for a set of elements using a 4th order Runge-Kutta method. We assume that the AAL region possesses a gas density of $n_{\textrm{H}} = 10^3~\textrm{cm}^{-3}$ and is illuminated by the `UV peak' SED at a distance of 2.3~kpc. Assuming a gas temperature $T = 10^4$~K, we then calculate recombination rate coefficients using the \citet{2006ApJS..167..334B} fits (assuming case A recombination), and collisional ionization rate coefficients using the data in \cite{1997ADNDT..65....1V}. We perform the integral in \cref{photoionization_rate} using the same photoionization cross-sections as used in \cloudy\ c13.00. Equilibrium values of $n_{M_i}$ are then calculated, assuming solar abundances. We assume these values to hold at a time $t = 0$. We next perturb the gas density such that $n(t > 0) = (1 + \delta) n(t = 0)$, where $-1 \leq \delta \leq \infty$, identical to the case of flux changes considered above. We do this for four different values of $(1 + \delta) = 0.1$, 0.01, 0.001 and 10. The latter is included for the interest of comparing both increasing and decreasing density. For example, if the AAL clouds are the result of shocked ISM clouds, we might expect them to be crushed prior to their subsequent expansion \cite[e.g.][]{FaucherGiguere:2012jk}. The resulting time-dependent evolution in the number density of \ion{H}{1}, \ion{O}{4}, \ion{O}{5}, \ion{O}{6}, \ion{Ne}{8} and \ion{Mg}{10} is shown in \Cref{non_equilibrium_figure} in \Cref{time_dependent_ionization_section}. Simple inspection of \Cref{non_equilibrium_figure} indicates that, for a decrease of more than an order of magnitude in gas density, with the exception of \ion{H}{1}, time-scales for reaching a new equilibrium are $> 100$~years. Comparing this to the expansion time-scale, we conclude that non-equilibrium effects are important if the density in the AAL region is dropping by orders of magnitude. Time-scales are orders of magnitude shorter for an increase in density, and so we expect that a decrease in density should form the dominant contribution to any non-equilibrium behaviour. We note that the situation described is not physical -- we expect a smooth change in density with time, not a step-function change. Nevertheless, lacking a physical motivation for the exact functional form, we present these results as an approximation. The non-equilibrium behaviour shown in \Cref{non_equilibrium_figure} indicates considerable variation in the rate of change of ionic number density, with the more highly ionized species changing more rapidly. Under photoionization equilibrium, the AAL region was found to possess a range in density and absorption path-length covering orders of magnitude. This situation was required due to the co-spatial existence of ionic species tracing gas with multiple ionization parameters. If ionic number densities are changing at highly variable rates, then is possible to envisage scenarios whereby the fractional abundances of ions spanning a large range in ionization potential can all be high. We therefore speculate that it may be possible to find non-equilibrium models that reproduce all of the observed column densities in a single phase, with a single density and absorption path length. Such a scenario may be desirable, as it is in better concordance with the near constant covering fraction seen across these ions. It is important to note that the calculations leading to inferred cloud lifetimes of $\lesssim 100$~years assume that the cloud sizes $l$ are approximately equal to the absorption path-length $l_{\textrm{abs}}$. However, in general, $l < l_{\textrm{abs}}$, which would make the cloud lifetimes even shorter. Scenarios where $l < l_{\textrm{abs}}$ have been put forward multiple times in the literature, typically to explain situations where the derived covering fractions in the data depend on velocity across the absorption profiles, and the ionization and/or true optical depth in the lines. This situation is referred to as inhomogenous partial coverage \cite[e.g.][]{deKool:2002by,Hamann:2004tu,Arav:2008fa}, where there are many small clouds having a power-law dependence in optical depth across their transverse extent. Covering fractions that vary in the way just described are not found in our data. Nevertheless, we cannot rule out the possibility that there are multiple small clouds along the line-of-sight. In such a scenario, the case for non-equilibrium evolution in the ionic number densities becomes more compelling. In summary, if the AAL region is not pressure supported, then explicit numerical calculations suggest that non-equilibrium effects may be important in these clouds. A cloud expansion time-scale of $\lesssim 100$ years is sufficiently short, that these effects might be confirmed with repeat observations. These observations will be crucial in determining appropriate non-equilibrium models for the data. \subsection{The connection to associated X-ray absorption} We next consider the link between associated \ac{uv} absorption and so-called `warm absorbers', often characterised by both bound-bound and bound-free absorption in X-rays. A key question is whether or not this absorption is predicted by the \ac{uv} absorption lines characterised here. In the most highly ionized gas, under photoionization equilibrium, the maximum predicted total hydrogen column density is $\log (N_{\textrm{H}} / \textrm{cm}^{-2}) \approx 20$, and the ionization parameter $\log U \approx 0.5$. We assume an upper limit on the gas metallicity of $[\textrm{O} / \textrm{H}] \approx 1$. Explicit photoionization calculations using these parameters predict the presence of bound-bound transitions, but no significant bound-free absorption in X-rays. The same result is obtained in \ac{cie} calculations. In this respect, and also in terms of their relatively small velocity shift from the \ac{qso}, the AALs in Q0209 are similar to those reported in e.g. \citet{Hamann:2000bi}. If there is continuous X-ray absorption, it should be in much more highly ionized gas with larger total column density. This gas may trace the bulk of an outflow that produces the gas condensations described in the previous section. However, in general, the gas giving rise to warm absorbers need not be co-spatial with \ac{uv} and/or optical \ac{aals}, especially since these absorption systems arise in gas with a wide range of physical parameters \cite[outflow velocities, ionization, covering fractions, distance from the \ac{qso} etc.; see for example][]{Ganguly:1999ia,Misawa:2007dd,Nestor:2008gb,Ganguly:2013hd,Muzahid:2013dm,Sharma:2013hk}. It is intriguing nonetheless, that the \ac{aals} in Q0209 show nearly identical velocity component structure over $\sim 300$~eV in ionization potential, suggesting that absorption lines from many ionization stages can indeed arise co-spatially. Simultaneous observations in the UV and X-rays will likely be required to gain deeper insights into the connection between warm absorbers and associated absorption lines in general \cite[e.g.][]{DiGesu:2013cf,Lee:2013jj}. \subsection{Outflow models} Before considering potential origins for the outflowing gas, we first perform a rough estimate of the mass and kinetic energy in the \ac{aals}. We assume the geometry of the outflowing gas traced by these data to be that of a thin, partially filled shell of material moving radially outwards from the centre of the \ac{qso}, the flux from which is modelled by the `UV peak' SED. Under this geometry, the mass depends on the distance from the \ac{qso} ($R \sim 2.3$~kpc), the total column density through the AAL region (we find a value $N_{\textrm{H}} \approx 2 \times 10^{20}~\textrm{cm}^{-2}$), and crucially, the global covering fraction, $\Omega$, of the AAL gas, as opposed to the line-of-sight covering fraction that we measure. A rough estimate of this quantity comes from the incidence rate of associated absorption systems like that seen in Q0209. \citet{Muzahid:2013dm} presented a sample of 20 quasar spectra observed with \ac{cos}, from which associated absorbers were selected based on the presence of \ion{Ne}{8} absorption. The incidence rate of these absorption systems was found to be 40\%, and we consider this to be the closest representative sample in the literature at present, although the general conclusions below are not sensitive to this value. We therefore express the total gas mass in the \ac{uv} AAL region as \begin{equation} M \approx 6 \times 10^7 \left(\frac{\Omega}{0.4}\right) \left(\frac{N_{\textrm{H}}}{2 \times 10^{20}~\textrm{cm}^{-2}}\right) \left(\frac{R}{2.3~\textrm{kpc}}\right)^2 M_{\odot}, \end{equation} where we have assumed a mean molecular weight per proton of $\mu_{\textrm{H}} = 1.4$. The total kinetic energy in this gas is then \begin{equation} K \approx 2 \times 10^{55} \left(\frac{M}{6 \times 10^7M_{\odot}}\right) \left(\frac{v}{200~\textrm{km s}^{-1}}\right)^2~\textrm{erg}. \end{equation} We can derive the average mass outflow rate $\dot{M}$, by dividing $M$ by the dynamical time-scale $R / v$, and subsequently derive the kinetic luminosity of the gas as $\dot{K} = 0.5\dot{M}v^2$. This gives values of $\dot{M} \approx 5 M_{\odot}~\textrm{yr}^{-1}$ and $\dot{K} \approx 7 \times 10^{40}~\textrm{erg s}^{-1}$. It is instructive to bear in mind that, while these quantities are useful, there are good reasons to believe that the gas clouds traced by the \ac{aals} may not travel a distance $R$, and are instead accelerated close to their observed location (see \Cref{structure_dynamics}). We consider two primary sources for the gas flow generating the \ac{aals}: (i) line-driven winds, and (ii) supernova-driven winds. Line-driven winds, initially accelerated through radiation pressure from the \ac{smbh} accretion disc, are commonly invoked to explain the winds traced by \ac{bals} and \ac{aals} with velocities of a few 1000 \kms, and are a major source of energy injection into the \ac{ism} in models of AGN feedback \cite[e.g.][]{Scannapieco:2004es,DiMatteo:2005hl,Hopkins:2010cf}. Specifically, these models require kinetic luminosities to be $\dot{K} \gtrsim 0.1$\% of the Eddington luminosity, $L_{\textrm{Edd}}$. For Q0209, $\log(L / L_{\textrm{Edd}}) \gtrsim 0$ (Done et al., in prep), which implies the kinetic luminosity in the AALs is at least two orders of magnitude below this level. In addition, models involving line-driven winds suggest they must be launched close to the \ac{smbh} at velocities of a few 100 \kms\ \cite[e.g.][]{Risaliti:2010jh}, which is already the velocity of the \ac{aals} seen here at much larger distances. If the \ac{aals} are pressure confined in a line driven wind such as this, they must encounter significant drag from a surrounding medium to slow them down, or keep them from accelerating to much larger velocities. In the more likely case that the \ac{aals} are formed in situ, a variety of velocities could in principle be observed. For example, in the radiative shock model of \cite{FaucherGiguere:2012jk}, AAL clouds will take a finite time to accelerate up to the velocity of the passing blast wave (see their equation (12)). Although the \ac{uv} AAL clouds in Q0209 contribute only a small percentage of the kinetic luminosity required for significant \ac{agn} feedback into the surrounding \ac{ism} (and \ac{igm}), a much larger percentage may be carried by an associated, much more highly ionized warm absorber, with higher total column density, detectable as bound-free absorption in X-rays \cite[e.g.][]{Crenshaw:2003hz,Gabel:2005jz,Arav:2007dx}. We note that bound-free X-ray absorption is by no means ubiquitous, with some warm absorbers now being detected via absorption \emph{lines} such as \ion{O}{7}. These can have $N_{\textrm{H}}$ consistent with that seen in associated UV absorption lines \cite[e.g.][]{DiGesu:2013cf}. Supernova driven winds are thought to drive fountains of gas a few kpc into the halos of massive galaxies, some of which is then expected to fall back on ballistic trajectories \cite[e.g.][]{Bregman:1980gn,Fraternali:2006eo,Fraternali:2008id,Marinacci:2010ho}. We find that the distance, velocity and density of the AAL gas in Q0209 is typical of galactic winds \cite[e.g.][]{Veilleux:2005ec,Creasey:2013gu}. The infalling velocity component $v_8$ also indicates that some of the AAL gas may be on a return trajectory back towards the disc of the \ac{qso} host. The time-scale derived in \cref{flow_time} is consistent with expected \ac{qso} lifetimes \cite[e.g.][]{Novak:2011gc}, so if the clouds are a result of supernova driven winds, it is possible these winds were launched during a starburst phase. If we assume that the starburst proceeded at $\sim 10 M_{\odot}~\textrm{yr}^{-1}$ and that this star formation results in one supernova per $100 M_{\odot}$ with energy $E_{\textrm{SN}} \sim 10^{51}~\textrm{erg}$, then $\dot{E} \sim 10^{42}~\textrm{erg s}^{-1}$. Only a fraction of this energy will be converted into the kinetic energy powering an outflow, and so the kinetic luminosity we derive for the \ac{aals} in Q0209 may be consistent with this simple estimation. At present, the data are consistent with both an outflow driven by AGN and starburst activity in Q0209. Future X-ray observations of this \ac{qso} may help to distinguish between these two possibilities. In particular, associated bound-free X-ray absorption tracing gas with high total column density and kinetic luminosity would favour an origin closely tied with the \ac{agn}.	14	3	1403.3338	We present HST/COS observations of highly ionized absorption lines associated with a radio-loud quasar (QSO) at z = 1.1319. The absorption system has multiple velocity components, with an overall width of ≈600 km s<SUP>-1</SUP>, tracing gas that is largely outflowing from the QSO at velocities of a few 100 km s<SUP>-1</SUP>. There is an unprecedented range in ionization, with detections of H I, N III, N IV, N V, O IV, O IV, O V, O VI, Ne VIII, Mg X, S V and Ar VIII. We estimate the total hydrogen number density from the column density ratio N(OIV) / N(OIV) to be log(n<SUB>H</SUB>/cm<SUP>-3</SUP>)∼3. Combined with constraints on the ionization parameter in the O IV bearing gas from photoionization equilibrium models, we derive a distance to the absorbing complex of 2.3≲R≲6.0kpc from the centre of the QSO. A range in ionization parameter, covering ∼two orders of magnitude, suggest absorption path lengths in the range 10<SUP>-4.5</SUP>≲l<SUB>abs</SUB>≲1pc. In addition, the absorbing gas only partially covers the background emission from the QSO continuum, which suggests clouds with transverse sizes l<SUB>trans</SUB>≲10<SUP>-2.5</SUP> pc. Widely differing absorption path lengths, combined with covering fractions less than unity across all ions pose a challenge to models involving simple cloud geometries in associated absorption systems. These issues may be mitigated by the presence of non-equilibrium effects, which can be important in small, dynamically unstable clouds, together with the possibility of multiple gas temperatures. The dynamics and expected lifetimes of the gas clouds suggest that they do not originate from close to the active galactic nuclei, but are instead formed close to their observed location. Their inferred distance, outflow velocities and gas densities are broadly consistent with scenarios involving gas entrainment or condensations in winds driven by either supernovae, or the supermassive black hole accretion disc. In the case of the latter, the present data most likely does not trace the bulk of the outflow by mass, which could instead manifest itself as an accompanying warm absorber, detectable in X-rays.	false	[ "Ar VIII", "gas densities", "multiple gas temperatures", "S V", "Ne VIII", "gas entrainment", "gas", "N V", "Mg X", "QSO", "N IV", "multiple velocity components", "simple cloud geometries", "s", "outflow velocities", "associated absorption systems", "absorption path lengths", "clouds", "velocities", "photoionization equilibrium models" ]	15.175407	7.574263	119
484123	[ "Barker, Adrian J.", "Dempsey, Adam M.", "Lithwick, Yoram" ]	2014ApJ...791...13B	[ "Theory and Simulations of Rotating Convection" ]	52	[ "Center for Interdisciplinary Exploration and Research in Astrophysics (CIERA) & Department of Physics and Astronomy, Northwestern University, 2145 Sheridan Road, Evanston, IL 60208, USA; Department of Applied Mathematics and Theoretical Physics, University of Cambridge, Centre for Mathematical Sciences, Wilberforce Road, Cambridge CB3 0WA, UK;", "Center for Interdisciplinary Exploration and Research in Astrophysics (CIERA) & Department of Physics and Astronomy, Northwestern University, 2145 Sheridan Road, Evanston, IL 60208, USA", "Center for Interdisciplinary Exploration and Research in Astrophysics (CIERA) & Department of Physics and Astronomy, Northwestern University, 2145 Sheridan Road, Evanston, IL 60208, USA" ]	[ "2014PhRvL.113y4501S", "2015ApJ...813..101V", "2015PEPI..246...52A", "2016A&A...592A..33M", "2016GeoJI.204.1120Y", "2016JFM...808..690G", "2016MNRAS.458.1548B", "2016MNRAS.459..939B", "2017ApJ...847L..16S", "2017LRSP...14....4B", "2018ApJ...856..132I", "2018JGRE..123..910L", "2018MNRAS.473.5002L", "2018sf2a.conf..135A", "2019ApJ...874...83A", "2019EAS....82....5M", "2019JFM...861R...5B", "2019arXiv190210594A", "2020ApJ...888L..31V", "2020ApJ...903...90A", "2020JFM...900R...1M", "2020MNRAS.491..923D", "2020MNRAS.493.5233C", "2020MNRAS.497.3400D", "2020MNRAS.497.4472V", "2020MNRAS.498.2270B", "2020MNRAS.498.3758J", "2020MNRAS.498.3782J", "2020PhRvR...2d3115A", "2020RvMP...92d1001S", "2020svos.conf..311A", "2021A&A...646A..48D", "2021JFM...912A..46L", "2021MNRAS.503.5789T", "2021PNAS..11822518V", "2022MNRAS.514.4111G", "2022PhFl...34j4117M", "2023A&A...673A...6D", "2023ApJ...950...73F", "2023ApJ...950L...4F", "2023ApJ...957L..23H", "2023Fluid...8..106H", "2023MNRAS.524.2661D", "2023SSRv..219...58K", "2023arXiv231004588W", "2024A&A...683A.221K", "2024ApJ...963...34W", "2024MNRAS.527.8245L", "2024MNRAS.528.6720L", "2024PhRvF...9d3501P", "2024SSRv..220...22F", "2024arXiv240604403V" ]	[ "astronomy", "physics" ]	6	[ "convection", "hydrodynamics", "stars: interiors", "stars: rotation", "turbulence", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Earth and Planetary Astrophysics", "Physics - Atmospheric and Oceanic Physics", "Physics - Fluid Dynamics" ]	[ "1954RSPSA.225..196M", "1958ZA.....46..108B", "1960ApJ...131..442S", "1961hhs..book.....C", "1962PhFl....5.1374K", "1966AnAp...29..489Z", "1971ARA&A...9..323S", "1974RSPSA.336...63H", "1977Icar...30..301G", "1979GApFD..12..139S", "1983JCoPh..49..241G", "1983JFM...126...75H", "1984ApJ...282..557H", "1984JCoPh..55..461G", "1990PhRvA..42.3650S", "1991ApJ...370..282C", "1994ApJ...421..245H", "1996ApJ...473..494B", "1998ApJ...493..955B", "2000ApJ...532..593M", "2000JFM...407...27G", "2001PEPI..128...51A", "2002ApJ...570..825B", "2002JFM...470..115C", "2003ARA&A..41..599T", "2003PhRvL..90c4502L", "2005A&A...438..403K", "2005A&A...444...25L", "2005ApJ...620..432R", "2005LRSP....2....1M", "2006ApJ...641..618M", "2006GeoJI.166...97C", "2006JFM...554..343G", "2009Icar..202..525K", "2009Icar..204..227J", "2009MNRAS.395.2056B", "2009MNRAS.400..176B", "2009Natur.457..301K", "2009PhRvE..80a5305S", "2009RvMP...81..503A", "2010MNRAS.404L..64L", "2010MNRAS.407.2451G", "2010NJPh...12g5022S", "2010NJPh...12k5002N", "2011A&A...531A.162K", "2011Icar..211.1258S", "2012Icar..219..428G", "2012JFM...691..568K", "2012PhRvL.109y4503J", "2013JFM...717..449K" ]	[ "10.1088/0004-637X/791/1/13", "10.48550/arXiv.1403.7207" ]	1403	1403.7207_arXiv.txt	\label{sec:intro} Rotating convection occurs in the interiors of many stars and planets. But there is no adequate theory for it yet, despite many decades of research. In order to determine the structure of a non-rotating star or planet, one typically employs mixing length theory (e.g.~\citealt{BohmVitense1958}), which despite its crudeness, accounts for the main structural feature of a convection zone: a nearly uniform entropy. But it remains unclear whether mixing length theory applies quantitatively, and how it may be extended to treat more subtle effects such as rotation and overshooting. Convection in astrophysical bodies---rotating or not---is difficult theoretically because the flow is turbulent. Compounding the difficulty, microscopic viscosities ($\nu$) and thermal diffusivities ($\kappa$) are typically extremely small in stars and planets, orders of magnitude smaller than accessible by experiment or simulation \citep[e.g.,][]{Spiegel1971,Miesch2005}. It is generally believed that bulk properties of turbulence should be independent of microscopic diffusivities in the limit that these are extremely small. But it remains unclear if that belief is correct, and if it is, whether experiments and simulations are adequately probing that limit. Rather than study the full problem in a star or planet, a common approach is to consider a simpler setup, Rayleigh-Benard convection (RBC), which has been extensively studied theoretically \citep[e.g.,][]{C1961,GrossmanLohse2000}, numerically, and experimentally (e.g.~\citealt{Ahlers2009}, and references therein). In RBC, a fluid layer is sandwiched between two horizontal plates. Convection is driven in the interior by holding the bottom plate hotter than the top. The fluid obeys the Boussinesq equations, in which the density is constant, the velocity field is incompressible, and temperature fluctuations give rise to vertical buoyancy forces \citep[e.g.][]{C1961}. This is perhaps the simplest set of equations that self-consistently evolve turbulent convection. Because of its simplicity, RBC is an ideal testbed for mixing length theory. The typical goal is to predict the heat flux through the fluid given the temperature drop between the plates. But the mixing length prediction, taken at face value, wildly disagrees with experimental and numerical results.\footnote{In the absence of rotation, mixing length theory predicts Nu $ = $ const$\times $Ra$^{1/2}$ \citep{Kraichnan1962}, where Ra (the Rayleigh number) quantifies the temperature drop and Nu (the Nusselt number) quantifies the heat flux. By contrast, experiments and simulations obtain an exponent $\approx 1/3$ at high Ra, i.e., at small $\nu$ and $\kappa$ \citep{Spiegel1971,ShraimanSiggia1990,Ahlers2009,2010NJPh...12k5002N}. The discrepancy increases as $\nu,\kappa\rightarrow 0$. Note that the mixing length prediction may be derived by dimensional analysis under the assumption that viscosity and thermal diffusivity play no role (\S \ref{mlt}).} The reason for the discrepancy is that a naive application of mixing length theory assumes a constant temperature gradient throughout the fluid. But in reality, most of the temperature drop occurs in extremely thin boundary layers near the top and bottom plates, where fluid velocities are nearly zero and heat is transported primarily by conduction. In fact, the relationship between temperature drop and heat flux can be accounted for by considering only the behavior of the boundary layers \citep{Malkus1954}. For {\it rotating} RBC, the situation is similar: \cite{King2012} extend the boundary layer analysis of \cite{Malkus1954} to include rotation, and thereby obtain impressive agreement with experiments and simulations (see also \citealt{King2009,King2013a}). However, for the purpose of explaining convection in astrophysical bodies, it is the behavior of the bulk of the fluid that is primarily of interest---not the boundary layers. The boundary layers in RBC are very different from the boundary of a convection zone in a star or planet. But one might expect that the turbulent dynamics of the interior fluid will be similar in the two cases. Therefore in this paper we focus on the dynamics of the interior fluid in an RBC-like system. We shall show that not only do the properties of the interior fluid converge in the limit of small diffusivities, but they converge to the prediction of the rotating mixing length theory first proposed by \cite{Stevenson1979}. The organization of this paper is as follows. We first set up the problem (\S \ref{Theory}), and then derive the predictions of mixing length theory (\S \ref{mlt}). Next, we test the theory in detail with a suite of numerical simulations that are similar to---but slightly different from---standard rotating RBC (\S \ref{Results}). We also run some comparison simulations with standard RBC (\S \ref{RBC}). We conclude with a summary and discussion (\S \ref{Conclusions}).	\label{Conclusions} We presented a simple derivation of mixing length theory in rapidly rotating convection, and then verified it with simulations. The theory, postulated by \cite{Stevenson1979}, predicts the properties of the convecting fluid under the assumption that they are independent of microscopic diffusivities ($\nu$ and $\kappa$). Equations \ref{eq:ml1}--\ref{eq:ml4} list the predictions for the mean temperature gradient, the velocity and temperature fluctuations, and the lengthscale of the modes that dominate heat transport. Our simulation results, summarized in Figure \ref{5}, agree remarkably well with the theory, across more than two orders of magnitude in rotation rate. We chose to focus on a very simple setup: Boussinesq convection in a box. But despite its simplicity, and despite the vast literature already devoted to the topic, the result remains under debate \citep[e.g.,][]{King2012,Julien2012}, largely because of the complicating effect of boundary layers. We circumvented this complication by focusing on the properties of the convecting fluid---i.e., between the boundary layers. We did this by fixing the flux, and examining the interior fluid's properties for increasingly small diffusivities. We thereby showed that the convecting fluid's properties converged to the prediction of mixing length theory as $\nu$, $\kappa\rightarrow 0$. Moreover, the numerical resolutions required to demonstrate convergence were relatively modest, after artificially thickening the boundary layers with heating/cooling zones. For example, our SNOOPY simulations had 256$^3$ gridpoints or fewer. Our numerical results provide strong support for those of \cite{Julien2012}, who simulate a set of reduced equations valid in the limit of rapid rotation. Our work lends confidence to mixing length theory's ability to accurately model highly turbulent convection. We hope to extend it to include a variety of more complicated---and realistic---effects, some of the most important of which are as follows: \begin{itemize} \item Including a background density gradient. One must then distinguish between entropy and temperature. The argument presented in \S \ref{mlt} should remain largely unchanged, after replacing temperature with entropy. But a possible complication is the asymmetry between upflows and downflows in the presence of a density gradient (e.g.~\citealt{Hurlburt1984,Cattaneo1991,Miesch2005}). \item Allowing for a more realistic geometry, i.e., quasi-spherical rather than cubical. The work presented in this paper strictly applies only to a small patch of the convective region near the poles of the star or planet. An intermediate step before considering the full spherical problem would be to allow rotation and gravity to be misaligned. \cite{Stevenson1979} predicts that in that case Eqs. \ref{eq:ml1}--\ref{eq:ml4} should be altered by replacing $\Omega\rightarrow\Omega\cos\theta$, where $\theta$ is co-latitude. But that has yet to be confirmed by simulations. \item Including the boundaries of a convection zone and the possibilities of penetration and overshooting into neighbouring stable layers (e.g.~\citealt{Hurlburt1994,Brummell2002,Rogers2005}) \item Allowing for the interaction between convection and differential rotation, and the generation of secondary flows. \item Including magnetic fields. \end{itemize} A large body of work has already been devoted to simulations of convection in rotating stars and planets, from Boussinesq (e.g.~\citealt{Hathaway1983,Tilgner2009,King2012}) to fully compressible \citep{Brummell1996,Brummell1998,Kapyla2005} Cartesian box simulations to Boussinesq \citep{Christensen2002,ChristensenAubert2006}, anelastic \citep{Glatz1984,Miesch2000,Kaspi2009,Jones2009,Gastine2012} and fully compressible simulations \citep{Kapyla2011} in spherical shell geometry. Given the evident complexities of some of these simulations, it is our view that a more complete understanding of simpler models is required to enable us to understand these simulation results. Our work complements the literature by definitively verifying the rotating mixing length theory described in \S \ref{mlt} for the case of Boussinesq convection in the polar regions of a planet or star. We anticipate that the theory described in this paper, as well as the extensions discussed above, will help provide a theoretical basis for the simulation results. Turning to astrophysical applications of the theory, we note first that rotation changes the entropy gradient relative to that predicted by standard (non-rotating) mixing length theory by an order-unity factor---at least for the Sun, where the rotation rate is comparable to the convective turnover time. Thus the inclusion of rotation will not substantially change static structure calculations, since it hardly affects the conclusion that convection zones have a near-constant entropy throughout \citep{Stevenson1979}. But a potentially important application is explaining the differential rotation profile of the Sun and other stars. In particular, small latitudinal entropy gradients drive differential rotation via the thermal wind equation (e.g.~\citealt{Thompson2003,Miesch2006,Balbus2009a,Balbus2009b}). Therefore to predict the differential rotation profile from first principles requires one to understand how the entropy gradient depends on latitude. It appears likely that rotating mixing length theory (at least when extended to the case in which rotation and gravity are misaligned) will provide an important piece towards solving this puzzle. Another potential application is to tidal dissipation in a convective star or planet that has an orbiting companion.\footnote{We thank Jeremy Goodman for pointing out this application to us.} This is important for understanding, for example, the tidal circularization of solar-type binary stars out to approximately ten day orbits. Previous work estimates the turbulent viscosity due to convection by employing non-rotating mixing length theory \citep{Zahn1966,GN1977}. It would be of interest to see how the predictions are affected by employing rotating mixing length theory.	14	3	1403.7207	We study thermal convection in a rotating fluid in order to better understand the properties of convection zones in rotating stars and planets. We first derive a mixing-length theory for rapidly rotating convection, arriving at the results of Stevenson via simple physical arguments. The theory predicts the properties of convection as a function of the imposed heat flux and rotation rate, independent of microscopic diffusivities. In particular, it predicts the mean temperature gradient, the rms velocity and temperature fluctuations, and the size of the eddies that dominate heat transport. We test all of these predictions with high resolution three-dimensional hydrodynamical simulations of Boussinesq convection in a Cartesian box. The results agree remarkably well with the theory across more than two orders of magnitude in rotation rate. For example, the temperature gradient is predicted to scale as the rotation rate to the four-fifths power at fixed flux, and the simulations yield 0.75 ± 0.06. We conclude that the mixing-length theory is a solid foundation for understanding the properties of convection zones in rotating stars and planets.	false	[ "rotation rate", "convection zones", "Boussinesq convection", "microscopic diffusivities", "fixed flux", "thermal convection", "convection", "heat transport", "simple physical arguments", "planets", "the imposed heat flux and rotation rate", "stars", "rapidly rotating convection", "the rotation rate", "order", "Cartesian", "Boussinesq", "Stevenson", "magnitude", "high resolution three-dimensional hydrodynamical simulations" ]	10.896348	14.430125	2
678409	[ "vom Marttens, R. F.", "Hipólito-Ricaldi, W. S.", "Zimdahl, W." ]	2014JCAP...08..004V	[ "Baryonic matter perturbations in decaying vacuum cosmology" ]	10	[ "Departamento de Física, Universidade Federal do Espírito Santo, Av. Fernando Ferrari, 514, Campus de Goiabeiras, CEP 29075-910, Vitória, Espírito Santo, Brazil", "Departamento de Ciências Naturais, Universidade Federal do Espírito Santo, CEUNES, Rodovia BR 101 Norte, km. 60, CEP 29932-540, São Mateus, Espírito Santo, Brazil", "Departamento de Física, Universidade Federal do Espírito Santo, Av. Fernando Ferrari, 514, Campus de Goiabeiras, CEP 29075-910, Vitória, Espírito Santo, Brazil" ]	[ "2014JCAP...10..060C", "2016JCAP...08..065H", "2016PhRvD..93f3503V", "2017IJMPD..2650019H", "2017PDU....15..114V", "2018EPJC...78..365H", "2019EPJC...79..355A", "2019JCAP...03..030C", "2019PhRvD..99d3521V", "2022EPJC...82..946A" ]	[ "astronomy", "physics" ]	4	[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Theory" ]	[ "1986ApJ...304...15B", "1986NCimB..93...36B", "1986PhLB..171..363O", "1987NuPhB.287..776O", "1987NuPhB.287..797F", "1990PhRvD..41..695C", "1991PhRvD..43.1075B", "1992PhRvD..46.2404C", "1994PhRvD..50.7725A", "1996ApJ...473...88R", "1996PhRvD..53.4280L", "1998PhRvD..58d3506O", "1999ApJ...517L...1O", "2000PhRvD..61f4007B", "2000PhRvD..61h3518M", "2000PhRvD..61h7304J", "2001ApJ...549..669H", "2001GReGr..33.1973V", "2001MPLA...16..633A", "2002IJMPD..11.1265M", "2002PhRvL..89h1302S", "2003GReGr..35..413Z", "2003IJMPA..18..811M", "2003PhLB..574..149S", "2003PhRvD..67d3502A", "2004IJMPD..13.1321C", "2004PhRvD..70f3529O", "2004PhRvD..70h7301H", "2005A&A...443..781G", "2005ApJ...633..560E", "2005CQGra..22..283W", "2005CQGra..22.3533B", "2005GReGr..37.1385B", "2005GrCo...11..277A", "2005JCAP...01..012S", "2005MNRAS.362..505C", "2005PhLB..611...27B", "2005PhLB..624..141W", "2005PhRvD..71j3504E", "2006IJMPD..15.1089S", "2006PhLB..637..357W", "2006PhRvD..73j3520B", "2006PhRvD..74l3507T", "2007A&A...466...11G", "2007ApJS..170..288H", "2007ApJS..170..377S", "2007CQGra..24..313M", "2008PhRvD..77d3513B", "2008PhRvD..77h3504C", "2009ApJ...700.1097H", "2009ApJS..182..543A", "2009JCAP...06..016H", "2009PhRvD..79f3527K", "2009PhRvL.103o1303P", "2010MNRAS.401.2148P", "2010PhRvD..82f3507H", "2011ApJS..192...18K", "2011JCAP...04..028Z", "2011JCAP...08..022P", "2011PhLB..695...41O", "2012ApJ...746...85S", "2012PhLB..716..165A", "2013ApJS..208...19H", "2013PhLB..727...37B", "2014A&A...571A..16P", "2014Natur.506..171P" ]	[ "10.1088/1475-7516/2014/08/004", "10.48550/arXiv.1403.0427" ]	1403	1403.0427_arXiv.txt	Explaining structure formation in the expanding Universe is one of the major topics in cosmology and astrophysics. According to the current main-stream understanding, dark matter (DM) and dark energy (DE) are the dynamically dominating components of the Universe \cite{wmap,wmap12,planck}. Baryons contribute only a small fraction of less than 5\% to the cosmic energy budget. The standard $\Lambda$CDM model does well in fitting most observational data but there is an ongoing interest in alternative models within and beyond General Relativity. A class of alternative models within General Relativity "dynamizes" the cosmological constant, resulting in so-called $\Lambda(t)$CDM models. Taking the cosmological principle for granted, cosmic structures represent inhomogeneities in the matter distribution on an otherwise spatially homogeneous and isotropic background. Dynamical DE models, $\Lambda(t)$CDM models are a subclass of them, have to deal with inhomogeneities of the DE component in addition to the matter inhomogeneities to which they are coupled. This makes these models technically more complex than the standard model. Ignoring perturbations of the DE component altogether may lead to inconsistencies and unreliable conclusions concerning the interpretation of observational data \cite{Park-Hwang}. Whether or not DE perturbations are relevant has to be decided on a case-by-case basis. The directly observed inhomogeneities are of baryonic nature. From the outset it is not clear that the inhomogeneities in the baryonic matter coincide with the inhomogeneities of the DM distribution. In particular, if DM interacts nongravitationally with DE, which happens in $\Lambda(t)$CDM models, while baryonic matter is in geodesic motion, this issue has to be clarified. A reliable description of the observed matter distribution has to consider the perturbation dynamics of the baryon fraction even though the latter only marginally influences the homogeneous and isotropic cosmic background dynamics. Then, in models with dynamical DE, the perturbations of baryonic matter will necessarily be coupled to the inhomogeneities of both DM and DE. In a general context, the importance of including the physics of the baryon component in the cosmic dynamics has been emphasized recently \cite{nature}. In this paper we extend a previously established decaying vacuum model \cite{humberto,julio,zimdahl,saulo,saulochap} by including a separately conserved baryon fluid with a four-velocity that differs from the four-velocity of the DM component. The basic ingredient of this model is a DE component with an energy density proportional to the Hubble rate. Moreover, it is characterized by an equation-of-state (EoS) parameter $-1$ for vacuum. Equivalently, the resulting dynamics can be understood as a scenario of DM particle production at a constant rate \cite{saulo} or as the dynamics of a non-adiabatic Chaplygin gas \cite{saulochap}. DE perturbations for this model are explicitly related to DM perturbations and their first derivative with respect to the scale factor in a scale-dependent way. It has been shown that on scales that are relevant for structure formation, DE fluctuations are smaller than the DM fluctuations by several orders of magnitude \cite{zimdahl}. Our analysis will be performed within a gauge-invariant formalism in terms of variables adapted to comoving observers \cite{VDF}. We shall derive a set of two second-order equations that couple the total fractional energy-density perturbations of the cosmic medium to the difference between these total perturbations and the fractional baryonic perturbations. The perturbations of the baryon fluid are then found as a suitable linear combination. As far as the background dynamics is concerned, our updated tests against observations from SNIa, BAO and the position of the first acoustic peak of the CMB spectrum confirm previous results \cite{Pigozzo2011}. Including the LSS data improves the concordance of the model compared with the case without a separately conserved baryon component. The joint analysis allows us to predict the baryon abundance of the Universe independently of the DM abundance. The corresponding probability density function (PDF) exhibits a pronounced peak at about 5\% for this abundance. This is a new feature which entirely relies on a separate consideration of the baryon fluid. The paper is organized as follows. In Sec.~\ref{model} we establish the basic relations of our three-component model of DE, DM and baryons. In Sec.~\ref{background} we recall the homogeneous and isotropic background dynamics of this model. Sec.~\ref{perturbations} is devoted to a gauge-invariant perturbation analysis which results in an explicit expression for the energy-density perturbations of the baryon fluid. In Sec.~\ref{observations} we test the model against observations using both background and LSS data. Our results are summarized in Sec.~\ref{conclusions}.	\label{conclusions} The components of the cosmological dark sector, DM and DE, are dominating the overall dynamics of the Universe. The small baryonic fraction of presently less than 5\% of the energy budget does only marginally influence the homogeneous and isotropic expansion history. With the help of data from SNIa, BAO and the position of the first peak of the CMB anisotropy spectrum we updated and confirmed previous results for the background. But as far as structure formation is concerned, the situation is different. The directly observed inhomogeneous matter distribution in the Universe is the distribution of visible, i.e., baryonic matter. While the standard scenario according to which the baryons after radiation decoupling are falling into the potential wells created by the DM inhomogeneities may suggest a similar distribution of DM and baryonic matter, the situation less clear if DM is in (non-gravitational) interaction with DE, while the (directly) observed baryon component is separately conserved. We have carried out a detailed gauge-invariant perturbation analysis for the baryon fluid in a $\Lambda(t)$CDM cosmology in which a cosmological term is decaying into DM linearly with the Hubble rate. Our key result is an expression for the fractional baryon energy-density perturbation for an observer comoving with the baryon fluid. Using the LSS data of the SDSS DR7 project we obtained the PDF for the baryon abundance of the Universe independently of the DM abundance. The best-fit value of this abundance is $\Omega_{B0} = 0.05 \pm 0.02$ (2$\sigma$ CL) in remarkable agreement with the result from primordial nucleosnthesis. A combined analysis, including also data from SNIa, BAO and CMB confirms this result. For the best-fit value of the DM abundance we found $\Omega_{M0}=0.32\pm0.02$ ($2\sigma$ CL) from the combined analysis (LSS+BAO+SNIa(Union2.1)+CMB) and $\Omega_{M0}= 0.35\pm 0.03$ ($2\sigma$ CL) from the LSS data alone. These values are higher than those for the standard model but smaller than the corresponding value for a $\Lambda(t)$CDM model without a separately conserved baryon component. Generally, the explicit inclusion of the baryon fluid improves the concordance between background and perturbation dynamics. Our results indicate that the investigated $\Lambda(t)$CDM cosmology, which does not have a $\Lambda$CDM limit, has a competitive background dynamics but as far as the baryon matter power spectrum is concerned, the $\Lambda$CDM model is clearly favored.	14	3	1403.0427	We consider the perturbation dynamics for the cosmic baryon fluid and determine the corresponding power spectrum for a Λ(t)CDM model in which a cosmological term decays into dark matter linearly with the Hubble rate. The model is tested by a joint analysis of data from supernovae of type Ia (SNIa) (Constitution and Union 2.1), baryonic acoustic oscillations (BAO), the position of the first peak of the anisotropy spectrum of the cosmic microwave background (CMB) and large-scale-structure (LSS) data (SDSS DR7). While the homogeneous and isotropic background dynamics is only marginally influenced by the baryons, there are modifications on the perturbative level if a separately conserved baryon fluid is included. Considering the present baryon fraction as a free parameter, we reproduce the observed abundance of the order of 5% independently of the dark-matter abundance which is of the order of 32% for this model. Generally, the concordance between background and perturbation dynamics is improved if baryons are explicitly taken into account.	false	[ "dark matter", "baryonic acoustic oscillations", "SDSS DR7", "baryons", "type Ia", "data", "first", "Hubble", "BAO", "the cosmic baryon fluid", "Union", "Constitution", "SNIa", "the cosmic microwave background", "account", "supernovae", "the dark-matter abundance", "CMB", "the anisotropy spectrum", "the Hubble rate" ]	11.150579	0.639166	89
483100	[ "Kaspi, Victoria M.", "Archibald, Robert F.", "Bhalerao, Varun", "Dufour, François", "Gotthelf, Eric V.", "An, Hongjun", "Bachetti, Matteo", "Beloborodov, Andrei M.", "Boggs, Steven E.", "Christensen, Finn E.", "Craig, William W.", "Grefenstette, Brian W.", "Hailey, Charles J.", "Harrison, Fiona A.", "Kennea, Jamie A.", "Kouveliotou, Chryssa", "Madsen, Kristin K.", "Mori, Kaya", "Markwardt, Craig B.", "Stern, Daniel", "Vogel, Julia K.", "Zhang, William W." ]	2014ApJ...786...84K	[ "Timing and Flux Evolution of the Galactic Center Magnetar SGR J1745-2900" ]	67	[ "Department of Physics, McGill University, Montreal, Quebec, H3A 2T8, Canada", "Department of Physics, McGill University, Montreal, Quebec, H3A 2T8, Canada", "Inter-University Center for Astronomy and Astrophysics, Post Bag 4, Ganeshkhind, Pune 411007, India", "Department of Physics, McGill University, Montreal, Quebec, H3A 2T8, Canada", "Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Department of Physics, McGill University, Montreal, Quebec, H3A 2T8, Canada", "Université de Toulouse, UPS-OMP, IRAP, Toulouse, France; CNRS, Institut de Recherche en Astrophysique et Planétologie, 9 Av. colonel Roche, BP 44346, F-31028 Toulouse Cedex 4, France", "Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Space Sciences Laboratory, University of California, Berkeley, CA 94720, USA", "DTU Space, National Space Institute, Technical University of Denmark, Elektrovej 327, DK-2800 Lyngby, Denmark", "Space Sciences Laboratory, University of California, Berkeley, CA 94720, USA; Lawrence Livermore National Laboratory, Livermore, CA 94550, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Department of Astronomy & Astrophysics, The Pennsylvania State University, 525 Davey Lab, University Park, PA 16802, USA", "Astrophysics Office, ZP 12, NASA Marshall Space Flight Center, Huntsville, AL 35812, USA", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA", "Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA 91109, USA", "Lawrence Livermore National Laboratory, Livermore, CA 94550, USA", "Goddard Space Flight Center, Greenbelt, MD 20771, USA" ]	[ "2014ApJ...786...62S", "2014ApJ...790...60A", "2014ApJ...793...88M", "2014MNRAS.442..372K", "2014arXiv1406.6458T", "2015AcASn..56S..73T", "2015ApJ...798...78C", "2015ApJ...798..120B", "2015ApJ...800...76S", "2015ApJ...800..109B", "2015ApJ...806..266L", "2015ApJ...807...93A", "2015ApJ...808...32T", "2015ApJ...808...81P", "2015ApJ...809..165Y", "2015ApJ...814....5Y", "2015ApJ...814...94M", "2015ApJ...815...15W", "2015JHEAp...7..137D", "2015MNRAS.446.1536P", "2015MNRAS.449.2685C", "2015MNRAS.451L..50T", "2015MNRAS.453..172P", "2015RAA....15.1467T", "2015RPPh...78k6901T", "2015SSRv..191..315M", "2015Sci...347.1103B", "2015arXiv151101901C", "2015aska.confE..39T", "2016A&A...589A.116M", "2016ApJ...818..122A", "2016ApJ...824..138Y", "2016ApJ...825..132H", "2016ApJ...833..189L", "2016JCAP...05..007M", "2016MNRAS.458.2088P", "2016MNRAS.461.2688P", "2016SCPMA..59a5752T", "2016arXiv160207738K", "2017AdSpR..59..736B", "2017ApJ...838...12Y", "2017ApJ...841...11T", "2017ApJ...847...85Y", "2017ApJ...851...17Y", "2017ApJS..231....8E", "2017MNRAS.471.1819C", "2017RAA....17...54C", "2017arXiv170509767T", "2018JASS...35..133P", "2018MNRAS.473.2304P", "2018MNRAS.474..961C", "2018smfu.book..321M", "2019AN....340..340H", "2019AN....340.1023G", "2019ApJ...871...73H", "2019ApJ...875..143W", "2019ApJ...877L..30G", "2019ApJ...882..173G", "2019BAAS...51c.292W", "2019MNRAS.483.4175Y", "2019MNRAS.485....2C", "2019MNRAS.487.1426B", "2020ApJ...894..159R", "2020ApJ...902....1H", "2021AstL...47..214K", "2023RAA....23h5005N", "2024FrASS..1094449A" ]	[ "astronomy" ]	10	[ "Galaxy: center", "pulsars: general", "stars: magnetic field", "stars: neutron", "X-rays: stars", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1984ApJ...276..325A", "1990Natur.348..707A", "1993ApJ...409..345A", "1996ApJ...465..487V", "2000ApJ...542..914W", "2000ApJ...543..340T", "2002ApJ...576..381W", "2003ApJ...588L..93K", "2003ApJ...596L..79K", "2004ApJ...605..378W", "2004ApJ...609L..67G", "2005SPIE.5898..360M", "2005SSRv..120..165B", "2007ApJ...663..497C", "2007ApJ...664L..35G", "2008ApJ...673.1044D", "2008ApJ...677..503T", "2008ApJ...679..681C", "2008ApJ...686..520Z", "2009ApJ...702..614D", "2009ApJ...703.1044B", "2009PASJ...61S..99H", "2010ApJ...719L.111Y", "2011ASSP...21..247R", "2011ApJ...739...94S", "2011MNRAS.414.1679E", "2012ApJ...748....3D", "2012ApJ...750L...6P", "2012ApJ...754L..12P", "2012ApJ...757...68A", "2012ApJ...761...66S", "2013ATel.5006....1D", "2013ATel.5009....1K", "2013ATel.5040....1E", "2013ATel.5073....1D", "2013ATel.5074....1D", "2013ATel.5124....1K", "2013ATel.5254....1K", "2013ApJ...763...82A", "2013ApJ...770..103H", "2013ApJ...770L..23M", "2013ApJ...770L..24K", "2013ApJ...774...92P", "2013ApJ...775L..34R", "2013MNRAS.429..688Y", "2013MNRAS.435L..29S", "2013Natur.497..591A", "2013Natur.501..391E", "2013arXiv1306.2264L", "2014ApJ...781..107K", "2014ApJ...784...37D" ]	[ "10.1088/0004-637X/786/2/84", "10.48550/arXiv.1403.5344" ]	1403	1403.5344_arXiv.txt	The recently identified Galactic Center magnetar SGR J1745$-$2900 has a brief but interesting observational history. It was discovered serendipitously during an ongoing monitoring program of the Galactic Center region with the Swift X-ray Telescope (XRT). On 2013 April 24 increased X-ray emission was detected from the SGR A* region \citep{drm+13}, followed the next day by a bright X-ray burst reported by Swift's Burst Alert Telescope (BAT) \citep{kkb+13,kbk+13}. {\it Swift} XRT observations that same day refined the position of the burster to within 2.8$''$ of Sgr A* \citep{kkb+13}. Target-of-Opportunity observations by the {\it Nuclear Spectroscopic Telescope Array (NuSTAR)} revealed 3.76-s pulsations from the new transient, and measured a spin-down rate that implies the presence of a neutron star having surface equatorial dipolar magnetic field strength\footnote{Estimated assuming simple magnetic braking in a vacuum via $B=3.2 \times 10^{19} \sqrt{P \dot{P}}$~G.} $1.6 \times 10^{14}$~G \citep{mgz+13}. This identified the source as a newly outbursting magnetar in the Galactic Center (GC) region. \citet{mgz+13} also showed that the source spectrum was well described by a blackbody of $kT=1$~keV plus a power law of index of 1.5. A {\it Chandra} observation later confirmed the GC association and localized the source to an offset from Sgr A* of only 2.4$''$ \citep{rep+13}. \citet{ekk+13} and \citet{sj13} reported on the detection of the radio pulsar counterpart, and \citet{efk+13} showed that the observed value of the rotation measure of the radio pulsar constrains the strength of the magnetic field near Sgr A*, which provides a unique test of radiative accretion theory for supermassive black holes. \citet{mgz+13} asserted that the spin-down rate of the magnetar is sufficiently large that bias due to dynamical effects in the GC region will be negligible, unless the measured spin-down rate were temporarily greatly enhanced \citep[e.g. due to glitch recovery; see][]{dkg08}. \citet{rep+13} argued dynamical effects may in principle be measurable at the $\sim$10\% level with long-term monitoring. However, the latter would have to be in spite of the likely continued fading of the source back to quiescence, as well as the often highly noisy nature of magnetar spin evolution post-outburst \citep[e.g.][]{wkg+02,gk04,dkg09,crj+08,dksg12}. Also of interest, independent of the Galactic Center location, is the magnetar outburst itself. Specifically, the flux and spectral evolution of magnetars post-outburst can potentially constrain the physics of neutron star magnetospheres and/or crustal and interior composition. In the former case, magnetar outbursts are hypothesized to be due to twists in localized magnetospheric regions of enhanced current known as ``j-bundles'' \citep{bel09}. Untwisting of j-bundles involves the return of current to a hot spot on the stellar surface, with gradually decreasing luminosity and temperature, predictions that can be tested by measuring flux and spectral evolution post-outburst. In the latter case, models of crustal cooling following a sudden heat injection can be fit to magnetar cooling curves, and can constrain e.g. the depth of the energy injection as well as the nature of the stellar temperature profile \citep[e.g.][]{kew+03,snl+12,pr12,akac13}. In either case, a significant hardness/flux correlation is expected and indeed thus far is generally observed \citep[e.g.][]{re11,sk11}. Here we report on continued {\it NuSTAR} and {\it Swift} XRT monitoring of SGR J1745$-$2900 post-outburst, specifically its timing and flux evolution. We show that the source's spin-down rate has recently undergone a large increase in magnitude, by over a factor of two. We suggest that the change in rate occurred coincidentally with a second X-ray burst seen by {\it Swift} BAT on MJD 56450 (7 June 2013) \citep{kbc+13}. If the burst association is correct, this change in spin-down rate occurred with no coincidental period glitch and without a large radiative change beyond the short $<$0.32-s burst and possibly slightly elevated flux on that day as reported by \citet{kbc+13}. We also report on the source's flux and spectral evolution $>$100 days post-outburst.	We have reported on X-ray observations made by {\it NuSTAR} and {\it Swift} over $\sim$120 days after the initial outburst of the Galactic Center magnetar SGR J1745$-$2900 in 2013 April. We find that the magnetar's spin-down torque has increased by a factor of nearly 3 compared with the spin-down rate initially measured by \citet{mgz+13}, with no evidence for any accompanying spin-up or spin-down glitch. We also show that the pulsar's 3--10~keV flux has declined monotonically by a factor of $\sim$2 over the first post-outburst 80 days, and that the blackbody temperature has decreased by $\sim$20\% over the initial 60 days, similar to what was reported by Rea et al. (2013), although we find evidence for a possible levelling-off of both flux and temperature. We observed a likely increase in the source's 10--30 keV flux 17 days post-outburst, but observe no accompanying timing or burst event. We find no evidence for the hardness/flux correlation commonly observed in magnetars, although this seems likely due to the narrow range of fluxes we have yet sampled. Further monitoring may yet reveal spectral softening as the source flux declines. We argue that the origin of the increase in the spin-down rate is likely to be magnetospheric, and that such torque variations, ubiquitous in magnetars, are likely to dominate over any timing signatures of motions related to the magnetar's proximity to Sgr A*. This work was supported under NASA Contract No. NNG08FD60C, and made use of data from the {\it NuSTAR} mission, a project led by the California Institute of Technology, managed by the Jet Propulsion Laboratory, and funded by the National Aeronautics and Space Administration. We thank the {\it NuSTAR} Operations, Software and Calibration teams for support with the execution and analysis of these observations. This research has made use of the {\it NuSTAR} Data Analysis Software (NuSTARDAS) jointly developed by the ASI Science Data Center (ASDC, Italy) and the California Institute of Technology (USA). We acknowledge the use of public data from the Swift data archive. This research has made use of the XRT Data Analysis Software (XRTDAS) developed under the responsibility of the ASI Science Data Center (ASDC), Italy. We thank the {\it Swift} SOT team for their work in scheduling. VMK receives support from an NSERC Discovery Grant and Accelerator Supplement, from the Centre de Recherche en Astrophysique du Qu\'ebec, an R. Howard Webster Foundation Fellowship from the Canadian Institute for Advanced Study, the Canada Research Chairs Program and the Lorne Trottier Chair in Astrophysics and Cosmology. RFA receives support from a Walter C. Sumner Memorial Fellowship. AMB was supported by NASA grants NNX-10-AI72G and NNX-13-AI34G. JAK was supported by supported by NASA contract NAS5-00136. JKV's work was performed under the auspices of the U.S. Department of Energy by Lawrence Livermore National Laboratory under Contract DE-AC52-07NA27344.	14	3	1403.5344	We present the X-ray timing and spectral evolution of the Galactic Center magnetar SGR J1745-2900 for the first ~4 months post-discovery using data obtained with the Nuclear Spectroscopic Telescope Array and Swift observatories. Our timing analysis reveals a large increase in the magnetar spin-down rate by a factor of 2.60 ± 0.07 over our data span. We further show that the change in spin evolution was likely coincident with a bright X-ray burst observed in 2013 June by Swift, and if so, there was no accompanying discontinuity in the frequency. We find that the source 3-10 keV flux has declined monotonically by a factor of ~2 over an 80 day period post-outburst accompanied by a ~20% decrease in the source's blackbody temperature, although there is evidence for both flux and kT having leveled off. We argue that the torque variations are likely to be magnetospheric in nature and will dominate over any dynamical signatures of orbital motion around Sgr A*.	false	[ "post-discovery using data", "Swift observatories", "Sgr A", "*", "orbital motion", "spin evolution", "SGR J1745", "the Nuclear Spectroscopic Telescope Array", "spectral evolution", "Galactic Center", "evidence", "kT", "post", "a bright X-ray burst", "~2", "a ~20% decrease", "the X-ray timing", "Swift", "SGR", "our data span" ]	5.3574	4.159593	69
746158	[ "Singh, Dinesh", "Wu, Kinwah", "Sarty, Gordon E." ]	2014MNRAS.441..800S	[ "Fast spinning pulsars as probes of massive black holes' gravity" ]	16	[ "Department of Physics, University of Regina, Regina, SK S4S 0A2, Canada; Department of Physics and Engineering Physics, University of Saskatchewan, Saskatoon, SK S7N 5E3, Canada", "Mullard Space Science Laboratory, University College London, Holmbury St Mary, Surrey RH5 6NT, UK", "Department of Physics and Engineering Physics, University of Saskatchewan, Saskatoon, SK S7N 5E3, Canada; Department of Psychology, University of Saskatchewan, Saskatoon, SK S7N 5E3, Canada" ]	[ "2014MNRAS.445.3415S", "2015PhLB..749..226R", "2016MNRAS.461.4295S", "2016PhRvD..94d4047P", "2017PhRvD..95f4017R", "2019MNRAS.485.1053L", "2019MNRAS.486..360K", "2019MNRAS.490.3262K", "2020A&A...644A.167K", "2020MNRAS.495..600K", "2020MNRAS.497.5421K", "2022JATIS...8a1013B", "2022MNRAS.511.3602L", "2022Univ....8...78W", "2023LRR....26....2A", "2024MNRAS.530.1118L" ]	[ "astronomy" ]	7	[ "black hole physics", "gravitation", "celestial mechanics", "pulsars: general", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1951RSPSA.209..248P", "1974RSPTA.277...59D", "1975ApJ...195L..51H", "1975ApJ...200L.131S", "1975Natur.256...23B", "1976ApJ...208L..55N", "1977ApJ...218L.109I", "1982Natur.300..728A", "1991ApJ...376...90L", "1991MNRAS.250..629K", "1996MNRAS.280..498D", "1996MNRAS.282.1038I", "1997ApJ...486L..27L", "1997MNRAS.284..576E", "1998ApJ...509..678G", "1999ApJ...514..388W", "1999MNRAS.308..863S", "2000ApJ...539L...9F", "2000ApJ...539L..13G", "2000ApJ...545..847M", "2002AJ....124.3270G", "2002ApJ...573..283P", "2002ApJ...578L..41G", "2002Natur.419..694S", "2003ApJ...582L..21B", "2003ApJ...591..891B", "2004A&A...424..733F", "2004AN....325..248P", "2004AnPhy.309..232E", "2004ApJ...615..253P", "2004CQGra..21.5441B", "2004MNRAS.351.1049M", "2004NewAR..48.1413C", "2005ASPC..328..147C", "2005PhLA..343....1C", "2005PhRvD..72h4033S", "2005Sci...307..892R", "2006ApJ...649...91F", "2006PhRvD..74l4006M", "2007A&A...474...55F", "2007ApJ...670...92G", "2007MNRAS.382.1922K", "2008ApJ...689.1044G", "2008LRR....11....8L", "2008Natur.455...78D", "2008PhRvD..78j4028S", "2009A&A...493.1161S", "2009A&A...507....1S", "2009ApJ...692.1075G", "2009ApJ...698..198G", "2009ApJ...705.1252W", "2010ApJ...708..834B", "2010ApJ...710.1063V", "2010ApJ...722...33S", "2010ApJ...722.1030W", "2011AIPC.1357..147F", "2011Ap&SS.336...67L", "2011ApJ...739...28X", "2011CQGra..28s5025P", "2011MNRAS.411.1575B", "2011MNRAS.412.2211G", "2012A&A...542A..44K", "2012A&A...545A..13Y", "2012ApJ...745...83A", "2012ApJ...745L..28A", "2012ApJ...747....1L", "2012GReGr..44..719I", "2012MNRAS.421.1517D", "2012Natur.491..729V", "2012Sci...338..355D", "2013A&A...555A..26L", "2013ApJ...770L..23M", "2013ApJ...770L..24K", "2013ApJ...778..145N", "2013MNRAS.430.1940R", "2014ApJ...783L...7D", "2014MNRAS.440L..86C" ]	[ "10.1093/mnras/stu614", "10.48550/arXiv.1403.7171" ]	1403	1403.7171.txt	\label{sec-intro} The presence of astrophysical black holes is inferred from various observations, such as the powerful electromagnetic radiation emitted by distant quasars. Although we have not yet ``seen'' black holes directly, %It is a still a great challenge to map the space-time around a black hole, which presumably exists, with certainty. it will soon be possible to image the massive central black hole in the Galaxy \citep[see][]{Doeleman08} and some nearby galaxies, e.g.\ M87, using submm VLBI observations \citep[see][]{Doeleman12,Dexter12,Asada12}. At present the strongest evidence for a massive black hole at the centre of the Galaxy comes from monitoring the motions of stars in the Sgr A* region. These observations have established that a large amount of unseen mass, $\approx 4.2 \times 10^{6}\, {\rm M}_{\odot}$ \citep{Ghez08, Gillessen09}, is enclosed within a volume having a radius $< 0.01$~pc at the Galactic Centre \citep[see][]{Eckart97, Ghez98}. The simple explanation for this unseen mass is a massive black hole \citep[see][]{Schoedel02, Ghez08, Gillessen09}. Naturally we would ask if this massive black hole is rapidly rotating or it is slowly rotating. Knowing the black hole's rotational rate has important astrophysical implications. It indicates how the black hole has evolved and perhaps how it was formed -- whether the black hole was built up by the merging of smaller black holes or simply by accreting a large amount of gas. A rotating black hole drags the space-time around it so stars and gas respond differently to Kerr and Schwarzschild black holes. X-ray spectroscopy of relativistic lines has been used to determine the spin of several black holes in active galactic nuclei \citep[e.g.\ for MCG-60-30-15,][]{Iwasawa96}. Theoretical calculations \citep[e.g.][]{Laor91, Kojima91} show that the profiles of relativistic emission lines, such as the Fe~K$\alpha$ line emitted from the surface of the accretion disk around a black hole, depend on the black hole's spin. However, in practice the reliability of the relativistic line method of black-hole spin measurement depends also on how well we model the accretion flow and on how well we understand the radiative processes that give rise to the diagnostic lines in the disk region close to the black hole event horizon \citep[see e.g.][]{Fuerst07, Svoboda09, Younsi12}. It is always a challenging task to measure the spin of a black hole much less to measure it with accuracy, be it a stellar-mass black hole in a binary system or a supermassive black hole in an active galactic nucleus. As shown in theoretical calculations, the parameter space is actually degenerate for the relativistic X-ray line profiles \citep{Laor91, Kojima91, Fuerst04}, thus one needs to resolve this issue properly to obtain a reliable black hole spin measurement. As for the black hole in the Galactic Centre, the lack of X-ray activity \citep{Baganoff03} in the current epoch implies an absence of an opaque gas accretion disk on which the relativistic X-ray lines are expected to form. Alternative methods for determining the black hole spin are therefore much needed. Observations have shown correlations between the mass of central black holes and the properties of the bulges of their host galaxies. In particular, a relatively tight $M$-$\sigma$ correlation is found for the nearby big galaxies \citep{Ferrarese00, Gebhardt00}, where $M$ is the mass of the central black hole and $\sigma$ is the velocity dispersion of the stars in the galactic bulge. For the Galaxy, the mass estimate of the central black hole and the measured velocity dispersion of the stars in the Galactic bulge are consistent with the empirical $M$-$\sigma$ relation derived for external galaxies \citep[see][]{Gueltekin09}. The most massive astrophysical black holes known to date have masses around $\sim 1.5 \times 10^{10}\, {\rm M}_{\odot}$ \citep[e.g. the central black hole in NGC~1277,][]{vandenBosch12}. Nuclear black holes with masses below $10^{6}\, {\rm M}_{\odot}$ in galaxies are not firmly established by stellar dynamics or by reverberation mapping \citep{Peterson04}, but there are observations indicating that some Seyfert galaxies may contain nuclear black holes in the mass range of $10^{5}-10^{6}\, {\rm M}_{\odot}$ \citep{Greene07, Xiao11}. The inclusion of small-bulge (low-mass) galaxies appears to steepen the slope of the $M$-$\sigma$ relation \citep{Graham11}. It is still unclear whether the least massive dwarf galaxies contain a central black hole (with $M_{\rm bh} \sim 10^{4}\, {\rm M}_{\odot}$) similar to the big elliptical galaxies. The lower mass limit for the central black holes in galaxies is not certain. Further extrapolation of the $M$-$\sigma$ relation to low-mass stellar spheroids implies that globular clusters would have nuclear black holes with mass $\sim 10^{3 }-10^{4}\, {\rm M}_{\odot}$ \citep[see][]{Lutzgendorf13}. There have been active searches for the intermediate-mass black holes (IMBH, black holes with masses $\sim 10^{2}-10^{4}{\rm M}_{\odot}$) in globular clusters as well as in dwarf galaxies. While there are claims of the discovery of IMBHs in globular clusters, there are also counterclaims of non-detection \citep[see e.g.\ the discussions in][]{vanderMarel10}. It is of great importance to accurately measure black hole masses in low-mass stellar spheroids and to properly resolve the issues regarding the low-end of the mass spectrum of non-stellar black holes. Here, in this work, we analyze the orbital motion of millisecond pulsars (ms-pulsars, fast spinning neutron stars) around a rotating black hole taking into account the effect of the pulsar's stress-energy tensor on the Kerr metric of the black hole. The compactness of neutron stars and the large mass ratios between nuclear black holes and the neutron stars allow a point-particle approximation for the neutron star, without compromising a proper treatment of the interaction between the spin of the neutron star and the black hole spin as manifested by the interaction between the spin of the neutron star and the the curvature of space-time induced by the black hole's gravity. Thus, the dynamics of these systems are well described by the Mathisson-Papapetrou-Dixon (MPD) equations for spinning test particles in an external space-time. We show how the orbital dynamics of a ms-pulsar is determined by spin-curvature coupling when it revolves around a black hole and how the orbital dynamics depend on the spin as well as the mass ratio between the black hole and the pulsar. In particular, we show that motion of the pulsar out of the usual orbital plane is substantial, relative to the orbital extent, if the mass of the rotating black hole is low enough. We organize the paper as follows. In \S 2 we present the formulation of the dynamics of systems containing a spinning neutron star orbiting around a massive black hole. In \S3 we give the scheme for solving the MPD equations and solutions of some example systems with parameters of astronomical interest. The significance of such binary systems and resulting astrophysical implications are discussed in \S4. Throughout in this work, unless otherwise stated, we use the natural unit system with $c =G =1$, where $c$ is the speed of light, and $G$ is the gravitational constant. We also adopt a signature of $+2$ for the space-time metric tensor.	Signals from pulsars orbiting in the strong field of moderate to massive black holes offer a means to determine the mass and spin of the central black hole. The non-linear nature of the general relativistic field equations $G_{\mu \nu} = 8 \pi T_{\mu \nu}$ means that the computation of the motion of anything more than a test particle in the strong field near a black hole generally requires numerical methods. In particular, the mass and spin of an orbiting neutron star will change the space-time geometry $G_{\mu \nu}$ through its stress-energy tensor $T_{\mu \nu}$. Until now, analysis of the motion of a pulsar near a black hole and that motion's effect on the observed pulsar signal have considered the motion of the pulsar as a test particle moving along a geodesic in the Kerr space-time of a rotating black hole. The effect of the mass and spin in the pulsar's stress-energy tensor on the pulsar's motion had not been previously considered. Here we have demonstrate that effect through the approximation given by the MPD equations. The MPD equations used here consider the effect of the first two moments of the pulsar's stress-energy tensor on the pulsar's motion. Our computations for the astrophysically important cases corresponding to intermediate mass black holes and the nuclear black holes of low-mass galactic spheroids show that the pulsar's spin leads to significant motion out of the usual orbital plane. The extent of the out of plane motion of a $1.5 {\rm M}_{\odot}$ pulsar becomes comparable to the extent of the orbit's radius for black holes of masses $10^{3}{\rm M}_{\odot}$. This motion therefore needs to be accounted for to properly interpret the timing of pulsar signals from a pulsar that is closely orbiting any intermediate mass black hole that may exist in globular cluster cores. Models of observed pulsar timing that use the MPD equations will therefore provide accurate measurements of masses and spins of central black holes in globular clusters and nuclear black holes in the galactic spheroids at the low end of the $M$-$\sigma$ relation. %\vspace*{2cm}	14	3	1403.7171	Dwarf galaxies and globular clusters may contain intermediate-mass black holes (10<SUP>3</SUP>-10<SUP>5</SUP> M<SUB>⊙</SUB>) in their cores. Estimates of ∼10<SUP>3</SUP> neutron stars in the central parsec of the Galaxy and similar numbers in small elliptical galaxies and globular clusters along with an estimated high probability of millisecond (ms)-pulsar formation in those environments have led many workers to propose the use of ms-pulsar timing to measure the mass and spin of intermediate-mass black holes. Models of pulsar motion around a rotating black hole generally assume geodesic motion of a `test' particle in the Kerr metric. These approaches account for well-known effects like de Sitter precession and the Lense-Thirring effect but they do not account for the non-linear effect of the pulsar's stress-energy tensor on the space-time metric. Here we model the motion of a pulsar near a black hole with the Mathisson-Papapetrou-Dixon (MPD) equations. Numerical integration of the MPD equations for black holes of masses 2 × 10<SUP>6</SUP>, 10<SUP>5</SUP> and 10<SUP>3</SUP> M<SUB>⊙</SUB> shows that the pulsar will not remain in an orbital plane with motion vertical to the plane being largest relative to the orbit's radial dimensions for the lower mass black holes. The pulsar's out-of-plane motion will lead to timing variations that are up to ∼ 10 μs different from those predicted by planar orbit models. Such variations might be detectable in long-term observations of ms pulsars. If pulsar signals are used to measure the mass and spin of intermediate-mass black holes on the basis of dynamical models of the received pulsar signal, then the out-of-plane motion of the pulsar should be part of that model.	false	[ "black holes", "pulsar motion", "geodesic motion", "planar orbit models", "pulsar signals", "intermediate-mass black holes", "motion", "dynamical models", "the lower mass black holes", "Models", "a rotating black hole", "plane", "ms pulsars", "a black hole", "planar orbit", "small elliptical galaxies", "globular clusters", "the received pulsar signal", "many workers", "ms-pulsar timing" ]	6.702193	3.983758	-1
746272	[ "Barton, Emma J.", "Chiu, Christopher", "Golpayegani, Shirin", "Yurchenko, Sergei N.", "Tennyson, Jonathan", "Frohman, Daniel J.", "Bernath, Peter F." ]	2014MNRAS.442.1821B	[ "ExoMol molecular line lists V: the ro-vibrational spectra of NaCl and KCl" ]	45	[ "Department of Physics and Astronomy, University College London, London WC1E 6BT, UK", "Department of Physics and Astronomy, University College London, London WC1E 6BT, UK", "Department of Physics and Astronomy, University College London, London WC1E 6BT, UK", "Department of Physics and Astronomy, University College London, London WC1E 6BT, UK", "Department of Physics and Astronomy, University College London, London WC1E 6BT, UK", "Department of Chemistry and Biochemistry, Old Dominion University, Norfolk 23529-0126, USA", "Department of Chemistry and Biochemistry, Old Dominion University, Norfolk 23529-0126, USA" ]	[ "2014JChPh.141n4312P", "2014MNRAS.445.1383Y", "2015ApJ...802..107W", "2015MNRAS.449.3613P", "2015MNRAS.451..634R", "2015MNRAS.454.1931P", "2015MolPh.113.1998L", "2015PCCP...1724598R", "2015PCCP...1724666M", "2016ApJ...825..150C", "2016IAUTA..29..137F", "2016JMoSp.327...73T", "2016JPhB...49j2001T", "2016MNRAS.456.4524Y", "2016PNAS..113E1424R", "2017JPhCS.810a2010Y", "2017JQSRT.203..511H", "2017MNRAS.468.1717M", "2017MNRAS.470..882W", "2017MolAs...8....1T", "2017RScI...88b3112L", "2018A&A...614A.131Y", "2018ASPC..515..137T", "2018ApJS..235....8B", "2018ApJS..237....8Y", "2018Atoms...6...26T", "2018JQSRT.210...44H", "2019A&A...627A..65C", "2019ApJ...872...54G", "2019PhyS...94l5402Z", "2020A&A...633A..39V", "2020A&A...636A..66P", "2020JQSRT.25507228T", "2021A&A...646A..21C", "2021A&A...655A..80D", "2021PCCP...2316390P", "2022A&A...658A..42P", "2022ApJ...940..129M", "2023A&A...678A..85C", "2023ApJ...942...66G", "2023Natur.614..664A", "2023PSST...32h5015O", "2023RemS...15..635P", "2024A&A...684A..25J", "2024arXiv240606347T" ]	[ "astronomy" ]	12	[ "molecular data", "opacity", "astronomical data bases: miscellaneous", "planets and satellites: atmospheres", "stars: low-mass", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1954PhRv...96..629H", "1957JChPh..27..573R", "1964PhRv..134R.863C", "1967JChPh..46.3749V", "1968JChPh..48.2824H", "1970JChPh..53..981D", "1981ApJS...45..621I", "1984ApJS...56..193S", "1987A&A...183L..10C", "1988CPL...147..133H", "1989JMoSp.134...98U", "1990CPL...169..599U", "1997JMoSp.183..360R", "2000JChPh.113.9766V", "2001MolPh..99.1535P", "2002ZNatA..57..663C", "2003Natur.421...45L", "2004JChPh.120.7939G", "2004JMoSt.695..243C", "2005JChPh.122j4103L", "2005JMoSt.742..215M", "2006MNRAS.368.1087B", "2007ApJ...668L.131M", "2011Natur.474..620P", "2012A&A...543A..48A", "2012ApJ...755...41S", "2012MNRAS.425...21T", "2012MNRAS.425...34Y", "2013AIPC.1545..186T", "2013ApJ...776...32M", "2013Icar..226.1673H", "2013JQSRT.121...78L", "2013JQSRT.130..284L", "2013MNRAS.434.1469B", "2014JMoSp.298....1T", "2014Natur.505...66K" ]	[ "10.1093/mnras/stu944", "10.48550/arXiv.1403.7952" ]	1403	1403.7952_arXiv.txt	NaCl and KCl are important astrophysical species as they are simple, stable molecules containing atoms of relatively high cosmic abundance. Na, K and Cl are the 15th, 20th and 19th most abundant elements in the interstellar medium \citep{04CaLeMu.KCl}. In fact NaCl could be as abundant as the widely-observed SiO molecule \citep{87CeGuxx.both}. \citet{87CeGuxx.both} reported the first detection of metal halides, NaCl, KCl, AlCl and, more tentatively, AlF, in the circumstellar envelope of carbon star IRC+10216. These observations have been followed up recently by \citet{12AgFoCe.both}, who also observed CS, SiO, SiS and NaCN. NaCl has also been detected in the circumstellar envelopes of oxygen stars IK Tauri and VY Caris Majoris \citep{07MiApWo.NaCl}. Another environment in which these molecules have been found is the tenuous atmosphere of Jupiter's moon Io. Submillimetre lines of NaCl, and more tentatively KCl, were detected by \citet{03LePaMo.NaCl} and \citet{13MoLeMo.KCl} respectively. NaCl has also been identified in the cryovolcanic plumes of Saturn's moon Enceladus alongside its constituents Na and Cl \citep{11PoScHi.NaCl}. K was also detected but the presence of KCl could not be confirmed. Furthermore NaCl and KCl are expected to be present in Super Earth atmospheres \citep{12ScLoFe.both} and may form in the observable atmosphere of the known object GJ1214b \citep{14KrBeDe.both}. The alkali chlorides are also of industrial importance as they are products of coal and straw combustion. Their presence in coal increases the rate of corrosion in coal fired power plants \citep{14YaKaLi.KCl}. Therefore it is important to monitor their concentrations, which can be done spectroscopically provided the appropriate data is available. The importance of NaCl and KCl spectra have motivated a number of laboratory studies, for example \citet{57RiKlxx.NaCl, 54HoMaSt.KCl, 88HoFuNa.NaCl, 89UeHoNa.NaCl, 64ClGoxx.both, 90UeHoKo.KCl}. The most recent and extensive research on KCl and NaCl spectra has been performed by \citet{97RaDuGu.NaCl}, whom investigated infrared emission lines of Na$^{35}$Cl, Na$^{37}$Cl and $^{39}$K$^{35}$Cl, \cite{04CaLeMu.KCl}, whom measured microwave and millimetre wave lines of $^{39}$K$^{35}$Cl, $^{39}$K$^{37}$Cl, $^{41}$K$^{35}$Cl, $^{41}$K$^{37}$Cl and $^{40}$K$^{35}$Cl, and \cite{02CaLeWi.NaCl}, whom recorded microwave and millimetre wave lines of Na$^{35}$Cl and Na$^{37}$Cl. Dipole moment measurements have been carried out by \citet{70LeWaDy.NaCl} for Na$^{35}$Cl and Na$^{37}$Cl, \citet{67WaDyxx.KCl} for $^{39}$K$^{35}$Cl and $^{39}$K$^{37}$Cl, and \citet{68HeLoMe.both} for $^{39}$K$^{35}$Cl, Na$^{35}$Cl and Na$^{37}$Cl. It appears the only theoretical transition line lists for these molecules are catalogued in the Cologne Database for Molecular Spectroscopy (CDMS), see \citet{cdms}. They were constructed using data reported in \citet{02CaLeWi.NaCl, 64ClGoxx.both, 89UeHoNa.NaCl, 70LeWaDy.NaCl} for NaCl, and \citet{04CaLeMu.KCl, 64ClGoxx.both, 67WaDyxx.KCl} for KCl. The lists are limited to $v$ = 4, $J$ = 159 and do not include a list for $^{41}$K$^{37}$Cl. In this paper we aim to compute more comprehensive line lists for the previously studied isotopologues and the first theoretical line list for $^{41}$K$^{37}$Cl. The ExoMol project aims to provide line lists on all the molecular transitions of importance in the atmospheres of planets. The aims, scope and methodology of the project have been summarised by \citet{jt528}. Lines lists for X~$^{2}\Sigma^{+}$ XH molecules, X = Be, Mg, Ca, and X~$^{1}\Sigma^{+}$ SiO have already been published (\citet{jt529, jt563} respectively). In this paper, we present ro-vibrational transition lists and associated spectra for two NaCl and four KCl isotopologues.	We present accurate but comprehensive line lists for the stable isotopologues of NaCl and KCl. Laboratory frequencies are reproduced to much more than sub-wavenumber accuracy. This accuracy should extend to all predicted transition frequencies up to at least $v$ = 8 and $v$ = 12 for NaCl and KCl respectively. New \ai\ dipole moments and Einstein $A$ coefficients are computed. Comparisons with the semi-empirical CDMS database suggest the intensities the pure rotational are accurate. The results are line lists for the rotation-vibration transitions within the ground states of Na$^{35}$Cl, Na$^{37}$Cl, $^{39}$K$^{35}$Cl, $^{39}$K$^{37}$Cl, $^{41}$K$^{35}$Cl and $^{41}$K$^{37}$Cl, which should be accurate for a range of temperatures up to at least 3000~K. The line lists can be downloaded from the CDS or from www.exomol.com. Finally, we note that HCl is likely to be the other main chlorine-bearing species in exoplanets. A comprehensive line lists for H$^{35}$Cl and H$^{37}$Cl have recently been provided by \citet{13LiGoHa.HCl} and \citet{13LiGoLe.HCl}.	14	3	1403.7952	Accurate rotation-vibration line lists for two molecules, NaCl and KCl, in their ground electronic states are presented. These line lists are suitable for temperatures relevant to exoplanetary atmospheres and cool stars (up to 3000 K). Isotopologues <SUP>23</SUP>Na<SUP>35</SUP>Cl, <SUP>23</SUP>Na<SUP>37</SUP>Cl, <SUP>39</SUP>K<SUP>35</SUP>Cl, <SUP>39</SUP>K<SUP>37</SUP>Cl, <SUP>41</SUP>K<SUP>35</SUP>Cl and <SUP>41</SUP>K<SUP>37</SUP>Cl are considered. Laboratory data were used to refine ab initio potential energy curves in order to compute accurate ro-vibrational energy levels. Einstein A coefficients are generated using newly determined ab initio dipole moment curves calculated using the CCSD(T) method. New Dunham Y<SUB>ij</SUB> constants for KCl are generated by a re-analysis of a published Fourier transform infrared emission spectra. Partition functions plus full line lists of ro-vibration transitions are made available in an electronic form as supplementary data to this paper and at www.exomol.com.	false	[ "<", "Cl", "ab initio potential energy curves", "accurate ro-vibrational energy levels", "ro-vibration transitions", "<SUP>41</SUP>K<SUP>37</SUP>Cl", ">K<SUP>37</SUP>Cl", "SUP>39</SUP", "supplementary data", "full line lists", "a published Fourier transform infrared emission spectra", "cool stars", "newly determined ab initio dipole moment curves", "Laboratory data", "exoplanetary atmospheres", "order", "KCl", "Accurate rotation-vibration line lists", "CCSD(T", "Fourier" ]	11.642955	12.248003	-1
800273	[ "Dent, James B.", "Krauss, Lawrence M.", "Mathur, Harsh" ]	2014PhLB..736..305D	[ "Killing the straw man: Does BICEP prove inflation at the GUT scale?" ]	15	[ "Department of Physics, University of Louisiana at Lafayette, Lafayette, LA 70504, USA", "Department of Physics and School of Earth and Space Exploration, Arizona State University, Tempe, AZ 85287, USA", "Department of Physics, Case Western Reserve University, Cleveland, OH 44106-7079, USA" ]	[ "2014IJMPD..2341001K", "2014JCAP...05..035A", "2014JCAP...07..018D", "2014JCAP...08..029D", "2014PhLB..733..276C", "2014PhLB..734...41G", "2014PhRvD..90b3506K", "2014PhRvD..90h3503D", "2014PhRvL.112q1302M", "2014PhyOJ...7...64K", "2015IJMPD..2441005D", "2015JCAP...04..011B", "2015PhRvD..91f3509G", "2016JCAP...02..023K", "2017JCAP...01..019D" ]	[ "astronomy", "physics" ]	9	[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]	[ "1981PhRvD..23..347G", "1982PhLB..108..389L", "1982PhLB..115..189R", "1983PhLB..125..445F", "1984NuPhB.244..541A", "1988ApJ...327..584S", "1992PhLB..284..229K", "1992PhRvL..69..869K", "1994PhRvD..49..692P", "1997PhRvL..79.1611P", "1997PhRvL..79.1615S", "1998PhRvD..58b3506T", "2008PhRvL.100m1302J", "2009JCAP...10..005F", "2010PhRvD..81l3504F", "2010PhRvD..82d4001K", "2011PhLB..695...26G", "2014A&A...571A..22P", "2014JCAP...08..029D", "2014PhRvD..89d7501K", "2014PhRvL.112q1301L", "2014PhRvL.112x1101B" ]	[ "10.1016/j.physletb.2014.07.046", "10.48550/arXiv.1403.5166" ]	1403	1403.5166_arXiv.txt			14	3	1403.5166	The surprisingly large value of r, the ratio of power in tensor to scalar density perturbations in the CMB reported by the BICEP2 Collaboration, if confirmed, provides strong evidence for Inflation at the GUT scale. While the Inflationary signal remains the best motivated source, a large value of r alone would still allow for the possibility that a comparable gravitational wave background might result from a self ordering scalar field (SOSF) transition that takes place later at somewhat lower energy. We find that even without detailed considerations of the predicted BICEP signature of such a transition, simple existing limits on the isocurvature contribution to CMB anisotropies would definitively rule out a contribution of more than 5% to r ≈ 0.2. We also present a general relation for the allowed fractional SOSF contribution to r as a function of the ultimate measured value of r. These results point strongly not only to an inflationary origin of the BICEP2 signal, if confirmed, but also to the fact that if the GUT scale is of order 10<SUP>16</SUP> GeV then either the GUT transition happens before Inflation or the Inflationary transition and the GUT transition must be one and the same.	false	[ "GUT", "scalar density perturbations", "scalar field", "the GUT transition", "simple existing limits", "strong evidence", "CMB anisotropies", "BICEP2", "the Inflationary transition", "the GUT scale", "the allowed fractional SOSF contribution", "Inflation", "≈", "the ultimate measured value", "r", "r.", "CMB", "place", "the isocurvature contribution", "the Inflationary signal" ]	11.46781	-0.589439	89
694294	[ "Biswas, Tirthabir", "Koivisto, Tomi", "Mazumdar, Anupam" ]	2014JHEP...08..116B	[ "Atick-Witten Hagedorn conjecture, near scale-invariant matter and blue-tilted gravity power spectrum" ]	20	[ "Department of Physics, Loyola University", "Nordita, KTH Royal Institute of Technology and Stockholm University, Roslagstullsbacken 23", "Consortium for Fundamental Physics, Physics Department, Lancaster University" ]	[ "2014CQGra..31o5010D", "2014JCAP...05..046K", "2014JCAP...07..033C", "2014JCAP...08..036M", "2014JCAP...10..075W", "2014MNRAS.442..670C", "2014PhLB..738..412G", "2014PhRvD..90l3524H", "2014arXiv1405.0513C", "2015IJMPD..2441005D", "2015MPLA...3040008C", "2015MPLA...3040009B", "2015PhRvD..92b3518C", "2015PhRvD..92l1304N", "2016PhRvD..93f3505B", "2016PhRvD..93j4001D", "2016PhRvD..94d3507N", "2023JHEAp..39...81V", "2023LRR....26....5A", "2024arXiv240613521B" ]	[ "astronomy", "physics" ]	3	[ "Cosmology of Theories beyond the SM", "String Field Theory", "High Energy Physics - Theory", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology" ]	[ "1987JETPL..45..709K", "1987PhRvD..35.3277S", "1988NuPhB.310..291A", "1989NuPhB.316..391B", "1989PhLB..220..125D", "1989PhRvD..40.2626D", "1992NuPhB.372..443T", "1993ppc..book.....P", "1995PhLB..361...45B", "2001PhRvD..64l3522K", "2003PhRvD..67d3518M", "2007JCAP...12..011B", "2007PhRvD..76b3510M", "2007PhRvL..98w1302B", "2009JCAP...06..037C", "2009LRR....12....2S", "2009PhRvD..80b3501A", "2009PhRvD..80b3519B", "2010JCAP...11..008B", "2010PhLB..684..177W", "2010PhRvD..82l3517B", "2010PhRvL.104b1601B", "2011PhR...497...85M", "2012JCAP...08..024B", "2012PhRvL.108c1101B", "2013ApJS..208...19H", "2013PhRvD..88b3517B", "2013PhRvD..88h3526B", "2014A&A...571A..16P", "2014CQGra..31o5010D", "2014PhRvD..90d7301G", "2014PhRvD..90f7301B", "2014PhRvD..90l3510A", "2014PhRvL.112x1101B", "2014SCPMA..57.1449W" ]	[ "10.1007/JHEP08(2014)116", "10.48550/arXiv.1403.7163" ]	1403	1403.7163_arXiv.txt			14	3	1403.7163	We will provide an interesting new mechanism to generate almost scale invariant seed density perturbations with a red spectrum, while keeping the gravitational wave spectrum blue-tilted in a stringy thermal contracting phase at temperatures beyond the Hagedorn temperature. This phase is often referred to as the Hagedorn phase where the free energy has been conjectured by Atick and Witten to grow more slowly than ordinary radiation. The primordial fluctuations are created by the statistical thermal fluctuations determined by the partition function, rather than quantum vacuum driven fluid dynamical fluctuations. Our mechanism assumes a non-singular bounce.	false	[ "quantum vacuum driven fluid dynamical fluctuations", "ordinary radiation", "temperatures", "quantum vacuum", "Hagedorn", "the statistical thermal fluctuations", "a stringy thermal contracting phase", "the Hagedorn temperature", "Witten", "Atick", "the Hagedorn phase", "The primordial fluctuations", "almost scale invariant seed density perturbations", "the partition function", "a non-singular bounce", "a red spectrum", "the free energy", "the gravitational wave", "This phase", "an interesting new mechanism" ]	10.726521	-0.567661	-1
769616	[ "Essam, A.", "Djurašević, G.", "Ahmed, N. M.", "Jurković, M." ]	2014NewA...32...16E	[ "A photometric study of the W UMa-type eclipsing binary system 1SWASP J160156.04 + 202821.6" ]	3	[ "National Research Institute of Astronomy and Geophysics, Department of Astronomy, Helwan, Cairo 11722, Egypt", "Astronomical Observatory of Belgrade, Volgina 7, 11060 Belgrade, Serbia", "National Research Institute of Astronomy and Geophysics, Department of Astronomy, Helwan, Cairo 11722, Egypt", "Astronomical Observatory of Belgrade, Volgina 7, 11060 Belgrade, Serbia" ]	[ "2015AJ....149..169J", "2016AJ....151..168K", "2018Ap&SS.363...15L" ]	[ "astronomy" ]	9	[ "Stars: binaries: close", "Stars: binaries: eclipsing", "Stars: individual: 1SWASP J160156.04 + 202821.6", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1956BAN....12..327K", "1967ZA.....65...89L", "1969AcA....19..245R", "1971ApJ...166..605W", "1979ApJ...231..502L", "1980MNRAS.193...79R", "1992Ap&SS.196..241D", "1992Ap&SS.197...17D", "1998A&AS..131...17D", "2000A&A...363.1081C", "2001ApJ...559..260W", "2004NewA....9..425D", "2005Ap&SS.296..221T", "2006PASP..118.1407P", "2011A&A...528A..90N", "2012A&A...542A.124L", "2012yCat.1322....0Z" ]	[ "10.1016/j.newast.2014.03.004", "10.48550/arXiv.1403.2746" ]	1403	1403.2746_arXiv.txt			14	3	1403.2746	B, V, R and I light curves of the eclipsing binary system 1SWASP J160156.04 + 202821.6 (J1601) have been constructed, for the first time, based on the CCD observations obtained using the 1.88 m telescope of Kottamia Astronomical Observatory (KAO), Egypt, during June, 2013. Twenty new times of minima (8 primary and 12 secondary) have been derived from the presented photometry, and a new ephemeris has been determined. The analysis of the corresponding light curves is made using Djurašević's inverse problem method. The light-curve asymmetry is explained using a Roche model with spots on system components. The analysis showed that the system is in an overcontact configuration (f<SUB>over</SUB> ∼ 13 %) with a small temperature difference between the components (ΔT =T<SUB>h</SUB> -T<SUB>c</SUB> ∼ 120K), indicating a good thermal contact. The analysis suggests that J1601 is an A-subtype of W UMa-type eclipsing systems, where the primary component is the hotter and more massive one.	false	[ "system components", "Kottamia Astronomical Observatory", "June", "ΔT =T<SUB", "1SWASP J160156.04", "W UMa-type eclipsing systems", "Egypt", "KAO", "a good thermal contact", "the eclipsing binary system", "J1601", "Djuraševićs inverse problem method", "CCD", "first", "a small temperature difference", "J160156.04", "the primary component", "curves", "1SWASP", "spots" ]	6.71964	11.055035	81
437685	[ "Castellano, M.", "Sommariva, V.", "Fontana, A.", "Pentericci, L.", "Santini, P.", "Grazian, A.", "Amorin, R.", "Donley, J. L.", "Dunlop, J. S.", "Ferguson, H. C.", "Fiore, F.", "Galametz, A.", "Giallongo, E.", "Guo, Y.", "Huang, K. -H.", "Koekemoer, A.", "Maiolino, R.", "McLure, R. J.", "Paris, D.", "Schaerer, D.", "Troncoso, P.", "Vanzella, E." ]	2014A&A...566A..19C	[ "Constraints on the star-formation rate of z ~ 3 LBGs with measured metallicity in the CANDELS GOODS-South field" ]	84	[ "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "Los Alamos National Laboratory, Los Alamos, NM, USA", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD, 21218, USA", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "UCO/Lick Observatory, Department of Astronomy and Astrophysics, University of California, Santa Cruz, CA, USA", "Department of Physics, University of California, Davis, CA, 95616, USA; Department of Physics and Astronomy, Johns Hopkins University, Baltimore, MD, 21218, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD, 21218, USA", "Cavendish Laboratory, University of Cambridge, 19 J.J. Thomson Avenue, Cambridge, CB3 0HE, UK; Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge, CB3 0HA, UK", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK", "INAF - Osservatorio Astronomico di Roma, via Frascati 33, 00040, Monteporzio, Italy", "Observatoire de Genève, Université de Genève, 51 Ch. des Maillettes, 1290, Versoix, Switzerland; CNRS, IRAP, 14 Avenue E. Belin, 31400, Toulouse, France", "Astronomisches Institut, Ruhr-Universität Bochum, Universitätsstraße 150, 44780, Bochum, Germany", "INAF - Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy" ]	[ "2014A&A...568L...8A", "2014A&A...569A..78V", "2014A&A...572A..41L", "2014ApJ...792..139T", "2014ApJ...796...95C", "2014Msngr.155...42F", "2015A&A...574A..19S", "2015A&A...575A..74S", "2015A&A...575A..96G", "2015A&A...576A.116V", "2015A&A...578A.105A", "2015A&A...582A..80T", "2015ApJ...799...32B", "2015ApJ...801...97S", "2015ApJ...803...34B", "2015ApJ...807..141P", "2015ApJ...814..161V", "2015MNRAS.451.2030D", "2015llg..book..463S", "2016A&A...587A.122A", "2016A&A...590A..31C", "2016A&A...591L...8S", "2016A&A...596A..75S", "2016ApJ...817...11H", "2016ApJ...818L...3C", "2016ApJ...820...73H", "2016ApJ...822...42O", "2016ApJ...831L...9N", "2016ApJ...833...37G", "2016ApJ...833...72B", "2016ApJ...833..178V", "2016ApJ...833..254S", "2016IAUS..319...45S", "2016MNRAS.459.1432P", "2016MNRAS.461.3886R", "2016MNRAS.462.3130M", "2017A&A...607A..30D", "2017A&A...608A.123D", "2017ApJ...838...29M", "2017ApJ...839...73C", "2017ApJ...847...76S", "2017ApJ...849...45C", "2017MNRAS.467.1360B", "2017MNRAS.467.4304V", "2017MNRAS.470.3006C", "2017MNRAS.471.4722M", "2017MNRAS.472..483F", "2017NatAs...1E..52A", "2018A&A...615A..77L", "2018ApJ...863..131F", "2018MNRAS.473...30C", "2018MNRAS.473.1258S", "2018MNRAS.473.2098M", "2018MNRAS.479...25M", "2019A&A...626A..45P", "2019A&A...630A.153A", "2019A&A...631A.123Y", "2019MNRAS.486..560S", "2019MNRAS.489...99B", "2019MNRAS.490.3309M", "2020A&A...637A..32B", "2020A&A...643A...4F", "2020MNRAS.495.1501C", "2020MNRAS.496..875R", "2021A&A...646A..39C", "2021A&A...649A..22M", "2021JCAP...01..010S", "2022A&A...659A..16L", "2022A&A...664A..73B", "2022ApJ...930..128R", "2022ApJ...938L..15C", "2022ApJ...940..135S", "2022MNRAS.511.4464A", "2022MNRAS.516.3532M", "2022PhRvD.105d3518S", "2023A&A...677A.138N", "2023ApJ...942L..27S", "2023ApJ...946...90M", "2023ApJ...948L..14C", "2023ApJ...948L..15V", "2023ApJ...952...84B", "2023MNRAS.526.5263B", "2024arXiv240211220N", "2024arXiv240520901M" ]	[ "astronomy" ]	6	[ "galaxies: distances and redshifts", "galaxies: evolution", "galaxies: high-redshift", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1955ApJ...121..161S", "1993ApJ...405..538B", "1994ApJ...429..582C", "1996A&AS..117..393B", "1998ARA&A..36..189K", "1998ApJ...497..618S", "1998ApJ...498..106M", "1999ApJ...521...64M", "2000ApJ...533..682C", "2001ApJ...554..803Y", "2001ApJ...562...95S", "2002ApJS..139..369G", "2003A&A...399...39R", "2003A&A...401.1063A", "2003ApJ...587..544F", "2003ApJ...592..728S", "2003MNRAS.344.1000B", "2004ApJ...615...98R", "2005ApJ...622..772C", "2005MNRAS.362.1054S", "2006A&A...454..423V", "2006AJ....132.1729B", "2007ApJ...670..928B", "2007MNRAS.376.1861F", "2007MNRAS.377.1024V", "2007PASP..119.1325L", "2008A&A...488..463M", "2009A&A...502..423S", "2009A&A...504..751S", "2009ApJ...692..778R", "2009ApJ...694.1517D", "2009ApJ...697.1493S", "2009ApJ...705..936B", "2009ApJ...706.1136O", "2009ApJS..183..244N", "2010A&A...524A..28C", "2010ApJ...709L..16O", "2010ApJ...712.1070R", "2010ApJ...725.1587B", "2010MNRAS.408.2115M", "2011A&A...532A..33G", "2011A&A...533A.119E", "2011ApJ...739L..40R", "2011ApJS..193...27W", "2011ApJS..197...35G", "2011ApJS..197...36K", "2011MNRAS.412.1828L", "2012A&A...537A..16F", "2012A&A...539A.136S", "2012A&A...540A..39C", "2012ApJ...744..154R", "2012ApJ...754...25R", "2012ApJ...754...83B", "2012ApJ...756..164F", "2012MNRAS.421.2568D", "2012MNRAS.423.1992S", "2013A&A...549A...4S", "2013A&A...549A..63K", "2013A&A...553A.132M", "2013ApJ...769...80A", "2013ApJ...772..136O", "2013ApJ...776...35K", "2013ApJS..206...10G", "2013ApJS..207...24G", "2013MNRAS.429..302C", "2013MNRAS.430.2885W", "2013MNRAS.432.2696M", "2013MNRAS.432.3520D", "2014A&A...563A..58T", "2014A&A...563A..81D", "2014ApJ...780..143A", "2014ApJ...783L..12V", "2014ApJ...793..115B" ]	[ "10.1051/0004-6361/201322704", "10.48550/arXiv.1403.0743" ]	1403	1403.0743_arXiv.txt	Lyman-break galaxies (LBG) represent by far the most numerous population of galaxies that we are able to observe in the early Universe. Their statistical distributions are becoming progressivley better and better constrained, most notably their UV luminosity function which has been determined from z$\simeq 3$ up to $z\simeq 8-9$ \citep[e.g.][]{Bouwens2007,Reddy2009,Castellano2010b,Grazian2011,McLure2013}. In turn, the debate about their physical properties and about how these properties change with redshift is more active. In particular, estimates of their dust extinction are needed to convert the UV luminosity density into a star-formation rate density (SFRD), and to constrain the amount of obscured star formation occurring in different systems at high redshift. An accurate determination of dust extinction in high-redshift galaxies is also fundamental to enable a proper comparison between predictions from galaxy evolution models and observations, so as to improve our understanding of the earliest stages of galaxy formation \citep[e.g.][]{Lacey2011,Somerville2012,Kimm2013}. To investigate the dust content of LBGs, great attention has been devoted to the study of the slope of the UV continuum \citep[e.g.][]{Castellano2012,Bouwens2012,Finkelstein2012,Dunlop2013}, which is mainly determined by dust absorption, but is also affected by other physical parameters, above all metallicity and age. Adopting the (reasonably accurate) assumption that the spectrum of LBGs between $\lambda \simeq 1500 \AA$ and $\lambda \simeq 2800 \AA$ can be represented by a simple power law $f_\lambda = \lambda^\beta$, these works have shown that high-redshift LBGs are generally found to have blue ($\beta$ $\sim$-2) slopes. Despite the remaining discrepancies among different works on the dependency of UV slopes on redshift and UV luminosity, these results have been univocally interpreted as an indication of relatively low dust obscuration by converting the observed $\beta$ into extinction assuming standard stellar populations. A thorough physical interpretation of these results remains, however, an open problem because of the intrinsic observational degeneracies. For a given extinction both lower metallicity stars and young ages are responsible for bluer UV slopes, while the contribution from nebular continuum produces a reddening of $\beta$. As recently noted by \citet{Wilkins2013} on the basis of galaxy formation models, any variation with redshift of the above factors can introduce systematic biases in the computation of dust extinction and of the corrected SFRD. Unfortunately, photometric data alone do not allow us to determine how different properties shape the observed $\beta$. In particular, while deep IR photometry leads to tighter constraints on the age of the stellar populations, stellar metallicity remains very poorly constrained even in the deepest photometric datasets. Similarily, it is still unclear where to place LBGs in a broad scenario of galaxy evolution, establishing the typical star-formation history (SFH) that led to their observed properties. Specifically, it would be important to establish whether the LBGs that we observe at different redshifts sample the same population observed at different epochs, while they assemble their stellar mass in a smooth, secular history of star formation, or if their SFH is more episode-driven and how it is related to possible phases of intense, dust-obscured star formation. In the latter case, individual galaxies might move in and out of the LBG selection criteria, and/or in different positions of the UV luminosity function over cosmic time. These questions can be, in principle, investigated by studying the spectral energy distribution (SED) of LBGs at various redshifts. However, several papers in the past have analyzed the SEDs of LBGs showing that, while stellar mass is reasonably well established, their age and SFH is more difficult to determine because of many degeneracies \citep[e.g.][]{Reddy2012,CurtisLake2013}. These degeneracies result from a set of uncertainties in the current observations of LBGs: photometry with high S/N is both scanty and difficult, especially in the crucial IR region where the contribution of previous generations of stars is appreciable; metallicity is generally not know, even for galaxies with spectroscopic redshift; and the large fraction of current LBG samples even lack spectroscopic redshifts, leading to additional uncertainties in the $k$--corrections and distance modulus. All these factors lead to larger uncertainties in the SED fitting of LBG samples that prevent us from constraining both their dust content and SFH. In this paper we take a different approach. Rather than selecting a complete sample of LBGs, we identify the small set of LBGs that have extraordinarily well-constrained properties, and perform a stringent, state-of-the-art SED fitting on them. In particular, we analyse a unique sample of galaxies at $z>2.9$ for which not only redshift but also metallicity (either stellar or gas-phase) has been measured, and exquisite deep photometry is available in all bands, from the optical to the crucial IR. To this purpose we have identified 14 galaxies at $2.9\lesssim z \lesssim 3.8$ in the GOODS-S field with measured metallicity (from deep spectroscopic surveys like AMAZE and GMASS) and we exploit the unique CANDELS dataset including observations from the U band to IRAC mid-IR, to perform SED fitting while fixing the metallicity of population synthesis models to the measured one. While the assumption of subsolar metallicities has recently become a common approach in the analysis of photometric samples at high redshift \citep[e.g.][]{Verma2007,Stark2009,Oesch2013,Alavi2014}, the objects considered here enable a detailed study of the effects of metallicity in the estimate of physical parameters. In particular, the availability of CANDELS WFC3 observations, of the deep K-band data of the HUGS-CANDELS survey (Fontana et al. in prep.), and of IRAC/SEDS data allows us to accurately sample the Balmer break at these redshifts to constrain the age of the objects in our sample. The available multi-wavelength data covering the rest-frame UV are also exploited to estimate extinction from the slope $\beta$ of the continuum under commonly adopted assumptions. Although our sample is not complete in a statistical sense, our galaxies are representative of relatively luminous LBGs, and we will show that the properties that we derive can provide useful information to settle the questions mentioned above. The plan of the paper is the following. In Sect.~\ref{sample} we present our sample and summarise the available spectroscopic and photometric information. In Sect.~\ref{properties} we discuss the SED-fitting estimates of their physical properties, obtained by varying assumptions on the SFH and on the contribution of nebular emission. In particular, we compare the resulting E(B-V) and SFRs to estimates obtained from UV slope and luminosity, and discuss independent constraints on the SFR from X-ray and FIR data. We exploit the results of our analysis to define a more appropriate $\beta-A_{1600}$ conversion equation which is then used to compute the SFRD at z$\sim$3 (Sect.~\ref{revisedformula}). A detailed discussion of the age of the galaxies in our sample is presented in Sect.~\ref{disc}. Finally, a summary is given in Sect.~\ref{summary}. Throughout the paper, observed and rest--frame magnitudes are in the AB system, and we adopt the $\Lambda$-CDM concordance model ($H_0=70km/s/Mpc$, $\Omega_M=0.3$, and $\Omega_{\Lambda}=0.7$).	\label{summary} In this work we have performed accurate SED fitting of a unique sample of 14 galaxies at $2.8\lesssim z \lesssim 3.8$ in the GOODS-South field. These galaxies constitute the only LBGs with both a spectroscopic measurement of their metallicity (either gas-phase or stellar, as measured from the AMAZE and GMASS surveys) and deep IR observations (obtained combining the CANDELS, HUGS, and SEDS surveys). Unfortunately, our galaxies do not make a complete sample in any statistical way; however, a posteriori they appear to have been selected from the general population of massive and luminous LBGs, and as such their analysis can shed some light on the general properties of the overall population of LBGs, or at least their brighter subsample. We have taken advantage of the 17-bands CANDELS catalogue to perform accurate SED fitting while fixing redshift and metallicity of population synthesis models to the measured values. For the spectral fit we use the BC03 models with Salpeter IMF, and we explore both different SFH (ranging from exponentially declining to rising) as well as models with or without the inclusion of nebular emission. Nebular emission has been computed both in the lines and continuum component following the procedure described in \citet{Schaerer2009}. We summarise here our findings, which are connected with two broad areas of investigation about high redshift LBGs: their dust content and their implied contribution to the global SFRD, and their ages and previous SFHs. \begin{itemize} \item \textbf{UV slope, dust content, and star--formation rates.} We first compared the SFR obtained through SED-fitting SFR$_{fit}$ with those estimated from the observed UV luminosity after correcting for the observed extinction (SFR$_{UV99}$), in the same manner as typically done on existing large LBG samples. We measured UV spectral slopes $\beta$ through a linear fit of HST magnitudes, and used the relevant extinction estimates \citep{Meurer1999} to estimate corrected SFR$_{UV99}$ according to the standard \citet{Madau1998} L$_{UV}$-SFR conversion. A comparison between SFR$_{fit}$ and SFR$_{UV99}$ shows that the latter are underestimated by a factor of 2-8, for all objects; SFR$_{fit}$ exceeds SFR$_{UV99}$ regardless of the assumed SFH, and the overestimate is larger for models without nebular emission (where it ranges typically between 3 and 8) rather than in models with nebular emission (ranging between 2 and 5). This result is supported by independent constraints on the radio (VLA), far-IR (Herschel) and X-ray (Chandra) emission of object CDFS-4417, which give SFR a factor of 2-4 larger than SFR$_{UV99}$ and in closer agreement with SFR$_{fit}$. This conclusion is also supported by the analysis of the far-IR stacking of the 13 sources that are not individually detected in Herschel data. The H$\beta$ luminosities measured for 11 of the objects also yield SFRs in agreement with SFR$_{fit}$, and confirm that the the objects are young and intensely star-forming. We demonstrate that these discrepancies are mostly due to the standard assumption of solar metallicity stellar populations underlying the widely used \citet{Meurer1999} UV slope-extinction conversion (Fig.~\ref{fig_zsun_age9}). On the basis of our results we deduce a new $\beta-A_{1600}$ relation, $A_{1600}=5.32+1.99\beta $ (Eq.~\ref{betaformula}), which is more appropriate for young subsolar metallicity LBGs. We note that this formula implies a dust-free UV slope as steep as $\beta=-2.67$, significantly bluer than the current assumption $\beta=-2.23$ based on the \citep{Meurer1999} formula. The $\beta-A_{1600}$ relation derived here is comparable to the one found by \citet{deBarros2012} and is consistent with theoretical predictions on the dust-free UV slope of high-z galaxies. Interestingly, Eq.~\ref{betaformula} also corresponds to an upward revision at z$\gtrsim$3 of the mean dust attenuation (L$_{IR}$/L$_{UV}$) vs. UV slope relation with respect to the M99 relation \citep[e.g.][]{Reddy2012b} by a factor of $\sim$2.5 for moderately extincted $\beta\gtrsim$-1.0 objects. It is interesting to explore the possible implications of these findings, under the assumption that these results can be extended to the overall LBG population. First, the common knowledge of negligible dust extinction in $\beta\lesssim -2.0 $ high-redshift galaxies would be seriously challenged, since this value for the UV slope can also be explained by the effect of extinction on sources with a steeper intrinsic spectrum \citep[see also][]{Dunlop2013}. This might bring the current theoretical predictions in better agreement with the observations, given that such models inevitably predict a rapid formation of substantial amounts of dust in high redshift LBGs \citep[e.g.][]{Lacey2011,Kimm2013}. Another important implication is on the contribution of LBGs to the global SFRD. We exploit our refined $\beta-A_{1600}$ relation, and use the average L$_{UV}$-SFR conversion for the objects in our sample to compute the z$\sim$3 SFRD on the basis of available estimates of the UV luminosity function and UV slope distribution at these redshifts. We find a dust corrected SFRD=$0.39~M_{\odot}/yr/Mpc^3$, more than two times higher than values based on old UV slope-extinction conversions. Adopting more conservative assumptions on the age of these subsolar metallicity galaxies, we anyway find SFRD estimates 40-60\% higher than those based on standard conversion equations. Of course, the effects on the SFRD discussed here are of comparable order to uncertainties between different IMFs. \item {\bf Ages and star-formation histories.} Finally we analysed in detail the constraints that our SED fitting is able to produce on the age and in general on the past evolutionary history of the objects in our sample (Sect.~\ref{disc}). We note that, on the basis of our fits, we are not able to constrain in a robust way the SFH. Both rising and declining models are found as best-fit solutions: 8 out of 14 LBGs are best-fit by exponentially declining models, the remaining are either fit with inverted-tau models or with rising-declining models having $age<\tau$, thus being in a rising SFH mode. In addition, the various SFHs are typically able to produce acceptable $\chi^2$ for most objects, such that the preference for a given SFH is not statistically robust even on individual objects. This is in agreement with the analysis of \citet{deBarros2012} of a large sample of $\sim$1700 LBGs. Our central result here is that we find very young best-fit ages for all our objects, in the range 10-500 Myr (Table~\ref{bestfits}). This result holds for any assumed SFH, both including or excluding nebular emission: in all these cases we find best-fit ages$\lesssim$500 Myr with very few exceptions. This finding is also supported by the measured L$_{H\beta}$/L$_{UV}$ of the 11 AMAZE sources in our sample, which is only consistent with an ongoing or a recently ($\sim$5 Myr) terminated burst. We have carefully explored whether this result is robust, given the expected prevalence of young stars in the overall SED, which may lead to important underestimates of the true age (the so-called overshining problem). We first decided to avoid the possibile complications of the SED fitting process and analysed a specific colour term designed to be particularly sensitive to age effects. We defined a ``Balmer colour'' (BCol=H160-0.5$\times$(K+3.6$\mu m$)) that brackets the 4000\AA~break, and computed BCol as a function of the $\beta$ value for templates of different SFHs (Fig.~\ref{fig_balmerplot}) at different ages. We show that six of the galaxies in our sample have a combination of low Balmer break ($\sim$0.0-0.5) and UV slope ($\beta\gtrsim$-2.5), unequivocally indicating very low ages (10-50Myr) for any SFH. The remaining galaxies have a position in the BCol-$\beta$ plane indicating Age$\lesssim$300Myr when declining or constant SFH are assumed, while only four objects with BCol$\sim$0.8 are compatible with exponentially increasing SFH templates of age$\sim$0.3-1.0 Gyr. Going back to the SED fitting technique, we also performed an estimate of the maximum age compatible for each SFH, defined as the largest age that produces a model with $P(\chi^2)>32$\%. The maximum age remains less than 0.5Gyr for at least half of the sample, regardless of the detailed SFH. The remaining 50\% can be reconciled with large ages (equivalent to formation redshifts around 10) only with increasing SFH. Similar results have been found using a specific set of models with double-burst SFH templates, where the relative intensity of the two bursts is left free. Even in this somewhat extreme case 9 out of 14 objects in the sample are found to have a very low fraction ($<$0.2) of old stellar population to their SED. However, significant uncertainties remain: constraints on the SFH are loose, while double-component libraries indicate a $>$50\% contribution from an old star-formation episode in four objects when constant $\Delta t_{burst}$=300Myr bursts are used, and for most of the objects when using shorter bursts. In addition, while our results are broadly consistent with a significant fraction of z$\sim$3 galaxies being very young, objects dominated by old stellar populations \citep[e.g.][]{Shapley2001} might be missing in such a small sample of bright LBGs. \end{itemize} The results summarised above show that tight constraints on metallicity and on the rest-frame optical regime are fundamental in order to shed light on two debated issues: 1) the estimate of dust-extinction and dust-corrected SFRs at high-redshift, and 2) the assembly history of Lyman-break galaxies. On the one hand, we have shown that taking into account the subsolar metallicity of stellar populations yields a significant revision of the UV slope-extinction conversion and of the corrected SFRD. On the other hand, the ages of a few 10-100Myr we find for our objects, the low amplitude of their Balmer-break, and the minor impact of older stars to their SEDs, suggest a particular scenario for the assembly of at least a sizeable fraction of the high-redshift LBGs. These luminous objects are most probably caught during huge star-formation bursts moving them on short timescales from the faint to the bright end of the luminosity function, rather than being the result of a constant, smooth evolution. The final word on these two questions will only come by the analysis of larger, and fainter samples. Extending the present work to a larger number of bright z$\gtrsim$3 LBGs will be possible thanks to intensive spectroscopic follow-up campaigns such as the ongoing VIMOS Ultra-Deep Survey, which is targeting 10000 galaxies including sky regions provided with deep near-IR observations. However, pushing this analysis to fainter galaxies, and to the redshifts approaching the reionisation epoch, is beyond the current limits of available instrumentation. Near infrared observations deeper than the ones presented here (m$_{lim}\sim$26.5 AB), and spectroscopic constraints of absorption features in $M>>M_{}$ LBGs, will be within reach only thanks to JWST and EELT.	14	3	1403.0743	<BR /> Aims: We aim to constrain the assembly history of high-redshift galaxies and the reliability of UV-based estimates of their physical parameters from an accurate analysis of a unique sample of z ~ 3 Lyman-break galaxies (LBGs). <BR /> Methods: We analyse 14 LBGs at z ~ 2.8-3.8 constituting the only sample where both a spectroscopic measurement of their metallicity and deep IR observations (CANDELS+HUGS survey) are available. Fixing the metallicity of population synthesis models to the observed values, we determine best-fit physical parameters under different assumptions about the star-formation history (SFH) and also consider the effect of nebular emission. For comparison, we determine the UV slope of the objects, and use it to estimate their SFR<SUB>UV99</SUB> by correcting the UV luminosity. <BR /> Results: A comparison between star-formation rate (SFR) obtained through SED-fitting (SFR<SUB>fit</SUB>) and the SFR<SUB>UV99</SUB> shows that the latter are underestimated by a factor of 2-10, regardless of the assumed SFH. Other SFR indicators (radio, far-IR, X-ray, recombination lines) coherently indicate SFRs a factor of 2-4 larger than SFR<SUB>UV99</SUB> and in closer agreement with SFR<SUB>fit</SUB>. This discrepancy is due to the solar metallicity implied by the usual β - A<SUB>1600</SUB> conversion factor. We propose a refined relation, appropriate for subsolar metallicity LBGs: A<SUB>1600</SUB> = 5.32 + 1.99 ∗ β. This relation reconciles the dust-corrected UV with the SED-fitting and the other SFR indicators. We show that the fact that z ~ 3 galaxies have subsolar metallicity implies an upward revision by a factor of ~1.5-2 of the global SFRD, depending on the assumptions about the age of the stellar populations. We find very young best-fit ages (10-500 Myr) for all our objects. From a careful examination of the uncertainties in the fit and the amplitude of the Balmer break we conclude that there is little evidence of the presence of old stellar population in at least half of the LBGs in our sample, suggesting that these objects are probably caught during a huge star-formation burst, rather than being the result of a smooth evolution. <P />Appendices are available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201322704/olm">http://www.aanda.org</A>	false	[ "Other SFR indicators", "SFR", "SUB", "nebular emission", "old stellar population", "subsolar metallicity LBGs", "UV99</SUB", "population synthesis models", "subsolar metallicity", "deep IR observations", "recombination lines", "SFH", "different assumptions", "-", "the other SFR indicators", "closer agreement", "SFR<SUB>UV99</SUB", "CANDELS+HUGS survey", "electronic form", "IR" ]	12.928271	6.842351	165
483368	[ "Caputi, K. I.", "Michałowski, M. J.", "Krips, M.", "Geach, J. E.", "Ashby, M. L. N.", "Huang, J. -S.", "Fazio, G. G.", "Koekemoer, A. M.", "Popping, G.", "Spaans, M.", "Castellano, M.", "Dunlop, J. S.", "Fontana, A.", "Santini, P." ]	2014ApJ...788..126C	[ "PdBI Cold Dust Imaging of Two Extremely Red H - [4.5] > 4 Galaxies Discovered with SEDS and CANDELS" ]	12	[ "Kapteyn Astronomical Institute, University of Groningen, P.O. Box 800, 9700 AV Groningen, The Netherlands", "SUPA, Institute for Astronomy, The University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK", "Institut de Radio Astronomie Millimétrique (IRAM), 300 rue de la Piscine, Domaine Universitaire, F-38406 Saint Martin d'Hères, France", "Centre for Astrophysics Research, Science & Technology Research Institute, University of Hertfordshire, Hatfield, AL10 9AB, UK", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Kapteyn Astronomical Institute, University of Groningen, P.O. Box 800, 9700 AV Groningen, The Netherlands", "Kapteyn Astronomical Institute, University of Groningen, P.O. Box 800, 9700 AV Groningen, The Netherlands", "INAF-Osservatorio Astronomico di Roma, Via Frascati 33, I-00040 Monteporzio, Italy", "SUPA, Institute for Astronomy, The University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK", "INAF-Osservatorio Astronomico di Roma, Via Frascati 33, I-00040 Monteporzio, Italy", "INAF-Osservatorio Astronomico di Roma, Via Frascati 33, I-00040 Monteporzio, Italy" ]	[ "2014IJMPD..2330015C", "2015ApJ...813...23V", "2016A&A...596A..75S", "2016ApJ...833..195W", "2017ApJ...835..286I", "2018ApJS..237...39A", "2020ApJ...901..168A", "2021A&A...649A.152K", "2021ApJ...908..146C", "2021ApJ...922..114S", "2023arXiv230515126K", "2023pcsf.conf...16K" ]	[ "astronomy" ]	9	[ "galaxies: active", "galaxies: high-redshift", "infrared: galaxies", "submillimeter: galaxies", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1955ApJ...121..161S", "1983ApJ...266..713O", "1998ApJ...498..541K", "1998ApJ...509..103S", "2000ASPC..217..299G", "2004A&A...424...23F", "2004ApJ...616...63Y", "2004ApJS..154...10F", "2004ApJS..154...25R", "2004ApJS..154..112L", "2005ApJ...622L.105H", "2005ApJ...631L..13D", "2006A&A...454..143C", "2006ApJ...651..142H", "2006MNRAS.366..609C", "2006MNRAS.367..349S", "2006MNRAS.372.1621C", "2007ApJ...670L..89W", "2007MNRAS.379.1599L", "2008ApJ...673L.127D", "2008ApJ...677..943D", "2008MNRAS.390.1117I", "2009ApJ...695L.176D", "2009ApJ...707.1201W", "2010A&A...514A..67M", "2010A&A...523A..27H", "2010ApJ...712..942M", "2010ApJ...717..379B", "2010ApJ...723..129K", "2010ApJ...724.1270T", "2010MNRAS.401..160A", "2011ApJ...742L..13H", "2011ApJS..197...35G", "2011ApJS..197...36K", "2011MNRAS.415.1479W", "2011Natur.470..233C", "2012ApJ...744..155W", "2012ApJ...750L..20C", "2012MNRAS.426.1845M", "2012Natur.486..233W", "2013ApJ...768..103C", "2013ApJ...769...80A", "2013MNRAS.432....2K", "2013MNRAS.432...53G", "2014ApJ...795....5O" ]	[ "10.1088/0004-637X/788/2/126", "10.48550/arXiv.1403.1435" ]	1403	1403.1435_arXiv.txt	\label{sec-intro} The powerful star formation and nuclear activity that led to the buildup of massive galaxies through cosmic time have been the subject of many studies. Most of these have focused on the cosmic time period elapsed between redshifts $z\sim 1.5$ and 3, when the cosmic star formation rate density had an overall peak \citep{hop06,beh10}, and the massive galaxy number density had a fast increase \citep[e.g.,][]{fon04,cap06a,sar06,kaj10}. At that time, stellar and nuclear activity were mostly obscured by dust, resulting in a high incidence of ultra-luminous infrared galaxies (ULIRGs). Indeed, a substantial fraction of the most massive galaxies were ULIRGs at $z\sim1.5-3$ \citep{dad05,cap06b}. The study of powerful star formation activity over the first few billion years of cosmic time ($z>3$) has proven to be more challenging, due to galaxy fainter fluxes, and the gradual decline of the cosmic star formation activity at high $z$. A notable exception to this challenge is offered by the study of bright sub-/millimetre selected galaxies, whose redshift distribution has a significant tail at $z>3$ \citep[e.g.,][]{war11,mic12}. However, the sensitivity limits of current sub-/millimetre surveys only allows for the study of the most extreme examples of early dust-obscured star formation, while a plausible population of more typical star-forming ULIRGs at $z>3$ is still to be found. An alternative approach for finding massive, dust-obscured starbursts at high $z$ consists of selecting bright mid-IR galaxies that are characterised by significantly red colours in their spectral energy distributions (SEDs). These red colours are the result of a redshifted 4000~$\rm \AA$ break and/or significant dust extinction. For example, different works have shown that optically faint, mid-IR bright galaxies are mostly dusty starbursts lying at $z \gsim 2$, and some also host active galactic nuclei (AGN) \citep[e.g.,][]{yan04,hou05,dey08}. Restricting this selection to those sources in which the significant flux drop occurs at near-IR wavelengths (observed $\lambda \approx 1-2 \, \rm \mu m$) should produce a redshift distribution biased towards even higher redshifts. Huang et al.~(2011) reported the existence of four galaxies selected with the {\em Spitzer Space Telescope} Infrared Array Camera \citep[IRAC;][]{faz04}, characterised by colours $H-[3.6]>4.5$ (AB). Their SED fitting suggests that these galaxies lie at $z\sim4-6$. Similarly, Wang et al.~(2012) analysed the SEDs of 76 IRAC galaxies with $K_s - [3.6]>1.6$ (AB), and found that about half of them are massive galaxies at $z\gsim3$. Making use of data from the {\em Spitzer} Extended Deep Survey \citep[SEDS;][]{ash13} and the {\em Hubble Space Telescope (HST)} Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey \citep[CANDELS;][]{gro11,koe11}, Caputi et al.~(2012; C12 hereafter) independently searched for these kinds of red galaxies in an area that is part of the UKIRT Infrared Deep Sky Survey \citep{law07} Ultra Deep Survey (UDS) field. C12 analysed the SEDs of 25 IRAC galaxies characterised by colours $H-[4.5]>4$ (AB), and concluded that between $\sim 45$ and $85\%$ of them are massive galaxies at $z>3$. Among the $z>3$ galaxies in C12, six have been detected by the Mid Infrared Photometer for {\em Spitzer} \citep[MIPS;][]{rie04} at $\tfm$, which at $z>3$ traces rest-frame wavelengths $\lambda_{\rm rest} < 6 \, \rm \mu m$, and thus indicates the presence of hot dust. For the brightest sources, this is likely due to the presence of an AGN. Understanding whether these galaxies are simultaneously undergoing a major episode of star formation requires us to follow them up at sub-/millimetre wavelengths, at which the cold-dust continuum emission can directly be probed. In this work we present PdBI 1.1~mm continuum observations towards the two brightest mid-IR galaxies in the $H-[4.5]>4$ sample analysed by C12. These interferometric observations have allowed us to achieve a spatial resolution of $\sim 1.8 \, \rm arcsec$ and sub-mJy sensitivities at millimetre wavelengths. Throughout this paper, all quoted magnitudes and colours are total and refer to the AB system \citep{oke83}. We adopt a cosmology with $\rm H_0=70 \,{\rm km \, s^{-1} Mpc^{-1}}$, $\rm \Omega_M=0.3$ and $\rm \Omega_\Lambda=0.7$. Stellar masses refer to a Salpeter (1955) initial mass function (IMF) over stellar masses $(0.1-100) \, \rm M_\odot$. \begin{figure} \epsscale{1.1} \plotone{stampcomp_pdbih45_27564.eps} \caption{Postage stamps of target id \#27564. From left to right: {\em HST} CANDELS f160w; SEDS IRAC $\ffm$; MIPS $\tfm$; PdBI clean 1.1~mm map. The shown field is of $\sim20\times 20$~arcsec$^2$ in all cases. \label{fig_maps27564}} \end{figure} \begin{deluxetable}{lcccccccc} \tabletypesize{\scriptsize} \tablecaption{Photometric properties of our two P\lowercase{d}BI targets. \label{table_targ}} \tablehead{\colhead{ID} & \colhead{RA (J2000)\tablenotemark{a}} & \colhead{DEC(J2000)\tablenotemark{a}} & \colhead{F160W} & \colhead{[4.5]} & \colhead{$S_\nu(\tfm)(\rm \mu Jy)$} & \colhead{$S_\nu(850 \, \rm \mu m)(mJy)$} & \colhead{$S_\nu(1.1 \rm mm)(mJy)$} } \startdata \#27564 & 02:17:16.35 & -05:14:43.1 & $24.89\pm0.05$ & $20.39\pm0.10$ & $599 \pm 13$ & $< 2.8 $ & $0.78 \pm 0.18$ \\ \#26857a & 02:17:51.69 & -05:15:07.2 & $24.39\pm0.14$ & $20.26\pm0.10$ & $334\pm 12$ & $< 2.8 $ & $<1.06$ \\ \#26857b & 02:17:51.62 & -05:15:03.6 & --- & --- & --- &$4.6 \pm 1.4$\tablenotemark{b} & $1.64 \pm 0.53$ \enddata \tablenotetext{a}{The RA and DEC values correspond to the IRAC coordinates, except for \#26857b, for which we quote the PdBI coordinates.} \tablenotetext{b}{The SCUBA2 $850 \, \rm \mu m$ source centroid is $\sim 3 \pm 4$~arcsec apart from our PdBI source centroid.} \end{deluxetable}	\label{sec-disc} Our PdBI detections towards two extremely red $H-[4.5]>4$ galaxies at $z>3$ are important for the following reasons. $\bullet$ Target \#27564 has a clear, $4.3 \sigma$-confidence millimetre counterpart, which confirms that this is a massive, AGN/star-forming composite galaxy at high redshifts. The millimetre flux density, which is completely dominated by star formation, indicates that this galaxy has an IR luminosity due to star formation $L_{\rm IR}^{\rm SFR} \approx 0.6-1.7 \times 10^{12} \, \rm L_\odot$, corresponding to $SFR\approx 200 \pm 100 \, \rm M_\odot/yr$. This implies that, from the star-formation point of view, this source is not like the typical hyper-luminous sub-/millimetre sources discovered thus far with single-dish millimetre telescopes at $z\sim 3-4$ \citep[e.g.,][]{mic10}. Rather, it is a modest ULIRG at $z>3$, such as those more commonly found by IR galaxy surveys at $z\sim 2-3$. Preliminary results on sub-millimetre number counts in deep ALMA observations \citep{kar13,ono14} suggest that, if the redshift distribution of faint sub-millimetre galaxies is comparable to that of the brighter sources currently known, then many more examples of these ordinary ULIRGs should be discovered at $z\sim 3-4$. $\bullet$ There is a tentative PdBI $3.1 \sigma$ detection at a distance of 3.7~arcsec of our target \#26857. The most likely scenario is that the two sources are unrelated, and the lack of another {\em Spitzer} or {\em HST} counterpart suggests that the PdBI detection corresponds to a new example of a very dusty starburst, like GN10, at high $z$. Our PdBI source also reveals that our $H-[4.5]>4$ galaxy is not the counterpart of the SCUBA2 detection in the same field, as a simple identification of the SCUBA2 source with the brightest IRAC source in the field would suggest. One could wonder whether the discovery of this new dusty source in the field of our target \#26857 is simply fortuitous. We believe that it likely is not: these red high-$z$ sources tend to be highly clustered \citep[e.g.,][]{tam10,cpk11}, so our finding of a GN10-like candidate source close to our $H-[4.5]>4$ target should perhaps not come as a surprise. However, this is not necessarily expected for all $H-[4.5]>4$ galaxies, as some of them do not have any close SCUBA2 detection (neither robust nor tentative). The analysis of the dusty IR SEDs of other $H-[4.5]>4$ galaxies suggests that not all of them are characterised by an important AGN component. At redshifts $z\sim3$, a pure dusty IR star-forming galaxy model is able to reproduce the mid-IR photometry and the sub-/millimetre photometric upper limits in all cases. Therefore, we conclude that our PdBI targets \#27564 and \#26857 could be prototypical of the brightest IRAC $H-[4.5]>4$ galaxies, but not all of them. As a general conclusion, we argue that the analysis of ultra-deep near and mid-IR data offers an alternative route to discover new sites of powerful star-formation activity over the first few billion years of cosmic time. We also conclude that associations between single-dish sub-millimetre sources and bright IRAC galaxies can be quite uncertain in some cases, and interferometric observations are necessary to study the dust-obscured star formation properties of the $H-[4.5]>4$ galaxies.	14	3	1403.1435	We report Plateau de Bure Interferometer (PdBI) 1.1 mm continuum imaging toward two extremely red H - [4.5] > 4 (AB) galaxies at z > 3, which we have previously discovered making use of Spitzer SEDS and Hubble Space Telescope CANDELS ultra-deep images of the Ultra Deep Survey field. One of our objects is detected on the PdBI map with a 4.3σ significance, corresponding to S_\nu (1.1 \, mm)=0.78 +/- 0.18 \, mJy. By combining this detection with the Spitzer 8 and 24 μm photometry for this source, and SCUBA2 flux density upper limits, we infer that this galaxy is a composite active galactic nucleus/star-forming system. The infrared (IR)-derived star formation rate is SFR ≈ 200 ± 100 M <SUB>⊙</SUB> yr<SUP>-1</SUP>, which implies that this galaxy is a higher-redshift analogue of the ordinary ultra-luminous infrared galaxies more commonly found at z ~ 2-3. In the field of the other target, we find a tentative 3.1σ detection on the PdBI 1.1 mm map, but 3.7 arcsec away of our target position, so it likely corresponds to a different object. In spite of the lower significance, the PdBI detection is supported by a close SCUBA2 3.3σ detection. No counterpart is found on either the deep SEDS or CANDELS maps, so, if real, the PdBI source could be similar in nature to the submillimeter source GN10. We conclude that the analysis of ultra-deep near- and mid-IR images offers an efficient, alternative route to discover new sites of powerful star formation activity at high redshifts.	false	[ "powerful star formation activity", "SCUBA2 flux density upper limits", "PdBI", "mid-IR images", "high redshifts", "CANDELS maps", "the ordinary ultra-luminous infrared galaxies", "Spitzer SEDS", "Hubble Space Telescope", "new sites", "SFR ≈", "SCUBA2", "gt", "a close SCUBA2 3.3σ detection", "mJy", "SEDS", "the PdBI detection", "the PdBI source", "Spitzer", "the PdBI map" ]	13.145721	6.865412	-1
483372	[ "Koshimoto, N.", "Udalski, A.", "Sumi, T.", "Bennett, D. P.", "Bond, I. A.", "Rattenbury, N.", "Abe, andF.", "Botzler, C. S.", "Freeman, M.", "Fukagawa, M.", "Fukui, A.", "Furusawa, K.", "Itow, Y.", "Ling, C. H.", "Masuda, K.", "Matsubara, Y.", "Muraki, Y.", "Ohnishi, K.", "Saito, To.", "Shibai, H.", "Sullivan, D. J.", "Suzuki, K.", "Suzuki, D.", "Sweatman, W. L.", "Takino, S.", "Tristram, P. J.", "Wada, K.", "Yock, P. C. M.", "MOA Collaboration", "Szymański, M. K.", "Kubiak, M.", "Soszyński, I.", "Pietrzynski, G.", "Poleski, R.", "Ulaczyk, K.", "Wyrzykowski, Ł.", "OGLE Collaboration" ]	2014ApJ...788..128K	[ "OGLE-2008-BLG-355Lb: A Massive Planet around a Late-type Star" ]	25	[ "Department of Earth and Space Science, Graduate School of Science, Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043, Japan; Microlensing Observations in Astrophysics (MOA) Collaboration.", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; Optical Gravitational Lensing Experiment (OGLE) Collaboration.", "Department of Earth and Space Science, Graduate School of Science, Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043, Japan; Microlensing Observations in Astrophysics (MOA) Collaboration.", "Department of Physics, University of Notre Dame, Notre Dame, IN 46556, USA; Microlensing Observations in Astrophysics (MOA) Collaboration.", "Institute of Information and Mathematical Sciences, Massey University, Private Bag 102-904, North Shore Mail Centre, Auckland, New Zealand; Microlensing Observations in Astrophysics (MOA) Collaboration.", "Department of Physics, University of Auckland, Private Bag 92019, Auckland, New Zealand", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, 464-8601, Japan", "Department of Physics, University of Auckland, Private Bag 92019, Auckland, New Zealand", "Department of Physics, University of Auckland, Private Bag 92019, Auckland, New Zealand", "Department of Earth and Space Science, Graduate School of Science, Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043, Japan", "Okayama Astrophysical Observatory, National Astronomical Observatory, 3037-5 Honjo, Kamogata, Asakuchi, Okayama 719-0232, Japan", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, 464-8601, Japan", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, 464-8601, Japan", "Institute of Information and Mathematical Sciences, Massey University, Private Bag 102-904, North Shore Mail Centre, Auckland, New Zealand", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, 464-8601, Japan", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, 464-8601, Japan", "Department of Physics, Konan University, Nishiokamoto 8-9-1, Kobe 658-8501, Japan", "Nagano National College of Technology, Nagano 381-8550, Japan", "Tokyo Metropolitan College of Industrial Technology, Tokyo 116-8523, Japan", "Department of Earth and Space Science, Graduate School of Science, Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043, Japan", "School of Chemical and Physical Sciences, Victoria University, Wellington, New Zealand", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, 464-8601, Japan", "Department of Earth and Space Science, Graduate School of Science, Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043, Japan", "Institute of Information and Mathematical Sciences, Massey University, Private Bag 102-904, North Shore Mail Centre, Auckland, New Zealand", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, 464-8601, Japan", "Mount John Observatory, P.O. Box 56, Lake Tekapo 8770, New Zealand", "Department of Earth and Space Science, Graduate School of Science, Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043, Japan", "Department of Physics, University of Auckland, Private Bag 92019, Auckland, New Zealand", "-", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; Universidad de Concepción, Departamento de Astronomia, Casilla 160-C, Concepción, Chile", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; Department of Astronomy, Ohio State University, 140 West 18th Avenue, Columbus, OH 43210,USA", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "-" ]	[ "2015ApJ...809...42C", "2016AJ....152..140B", "2016ApJ...820....4S", "2016ApJ...824..139H", "2016ApJ...830..150P", "2016ApJ...833..145S", "2016MNRAS.457.1320T", "2017AJ....153....1K", "2017AJ....153..143M", "2017AJ....154....1H", "2017AJ....154....3K", "2017AJ....154...59B", "2017AJ....154..103B", "2018AJ....155..263S", "2018AJ....156..113B", "2018AJ....156..236Z", "2018arXiv180410136N", "2018exha.book.....P", "2019MNRAS.487.4603T", "2020AJ....159...58R", "2020AJ....159...68B", "2020AJ....159..268K", "2020arXiv200902329B", "2021AJ....162...17K", "2021ApJ...917...78K" ]	[ "astronomy" ]	7	[ "gravitational lensing: micro", "planetary systems", "Astrophysics - Earth and Planetary Astrophysics" ]	[ "1972epcf.book.....S", "1985prpl.conf.1100H", "1988PASP..100.1134B", "1991ApJ...374L..37M", "1992ApJ...396..104G", "1993ARA&A..31..129L", "1993PASP..105.1342S", "1993prpl.conf.1031W", "1994AcA....44..227U", "1995ApJ...454L.125A", "1995Natur.378..355M", "1996ApJ...472..660B", "1996Natur.380..606L", "1997ApJ...480..647D", "2000A&A...363.1081C", "2000ApJ...542..785G", "2001MNRAS.327..868B", "2003AcA....53..291U", "2003ApJ...591..204S", "2003ApJ...592..172H", "2003ApJS..148..195V", "2004A&A...426..297K", "2004ApJ...604..388I", "2004ApJ...606L.155B", "2004ApJ...612L..73L", "2005ApJ...626.1045I", "2005ApJ...628L.109U", "2006ApJ...644L..37G", "2006ApJ...646..505B", "2006ApJ...647L.171B", "2006ApJ...649..436E", "2006Natur.439..437B", "2007ApJ...660..781B", "2007ApJ...670..833J", "2008ApJ...673..502K", "2008ApJ...684..663B", "2008ExA....22...51S", "2008PASP..120..531C", "2008Sci...319..927G", "2008Sci...322.1348M", "2009A&A...497..497G", "2009ApJ...695..970D", "2009ApJ...698.1826D", "2010AcA....60..197J", "2010ApJ...710.1641S", "2010ApJ...710.1800G", "2010ApJ...713..837B", "2010ApJ...716.1408B", "2010ApJ...720.1073G", "2010PASP..122..905J", "2010Sci...330..653H", "2011A&A...529A.102B", "2011AcA....61...83S", "2011ApJ...736...19B", "2011ApJ...741...22M", "2011Natur.473..349S", "2012A&A...540A..78K", "2012ARA&A..50..411G", "2012ApJ...755..102Y", "2012ApJ...757..119B", "2012Natur.481..167C", "2013A&A...549A.109B", "2013A&A...549A.147B", "2013A&A...552A..70K", "2013ApJ...762L..28H", "2013ApJ...768..129C", "2013ApJ...769...88N", "2013ApJ...778...38H", "2013ApJ...779...91F", "2014ApJ...780...54B", "2014ApJ...780..123S", "2014ApJ...781...28M", "2014ApJ...782...47P", "2014ApJ...782...48T", "2014ApJ...788...73Z", "2014MNRAS.439..604S" ]	[ "10.1088/0004-637X/788/2/128", "10.48550/arXiv.1403.7005" ]	1403	1403.7005.txt	Until the first detection of an exoplanet in 1995, planet formation theories referred to the formation of the Solar System. The standard core accretion model \citep{saf72,hay85,lis93} was believed to be fairly well established, although some problems such as the formation of planetesimals \citep[e.g.]{wei93,dom97} remained. According to this theory, gas giants, such as Jupiter or Saturn, are formed slightly outside the "snow line" where the protoplanetary disk becomes cold enough for water to condense. However, this theory did not predict the discovery of ``Hot Jupiters" \citep{may95}, which are planets of about a Jupiter mass with orbits lie far inside that of Mercury. Since then, over 1000 exoplanets (and over 3500 candidates) have been detected. The core accretion model now includes the possibility of migration \citep{lin96} to explain Hot Jupiters, but it still has difficulty which is derived from the standard model for the origin of the solar system is a most generally accepted planet formation scenario, but it can't explain all the forms of the exoplanets. For example, the theoretical prediction of a paucity of the planets with masses of 10 - 100 $M_{\oplus}$ in short period orbits \citep{ida04} is inconsistent with the results from radial velocity studies \citep{how10}. Moreover, today's core accretion model predicts few gas giants orbiting red dwarf at any separation \citep{lau04,ken08}, and this is confirmed by observations from radial velocity for massive gas giants orbiting inside the snow line \citep{end06,joh07,cum08,joh10}. Early statistical results from the gravitational microlensing method \citep{gou10b,sum10,cas12} indicate that low-mass, Saturn-like, gas giants are more common around low-mass stars than Jupiters beyond the snow line, and this is confirmed by radial velocity observations \citep{montet14}. But, the gravitational microlensing method \citep{mao91,gou92}, has also revealed several super-Jupiter mass planets orbiting just outside of the snow line of their late type host stars \citep{ben06,gau08,don09a,don09b,bat11,kai13,tsa13,shv13} although a quantitative analysis of planetary frequency as a function of host star mass has not yet been completed. The gravitational microlensing method is capable of discoveries of planets with mass down to the Earth mass just outside of the "snow-line" \citep{ben96}. In terms of the sensitivity region, it is very important for planetary formation theory that the microlensing method is complementary with the other methods, the radial velocity method \citep{but06,bon11} and the transit method \citep{bor11}, which are sensitive to close and relatively massive planets, and the direct imaging method \citep{mar08} which has sensitivity to giant planets with orbital semi-major axes greater than several dozen AU. Also, because microlensing does not rely upon any light from the host star or planet of the lens system for detection \citep{gau12}, it is possible to detect a planet around a star which is too faint to detect by the other methods \citep{ben08} or even a planetary mass object which belongs to no host star \citep{sum11}. The microlensing events that are searched for planetary signals are discovered by two microlensing survey groups, the Microlensing Observations in Astrophysics group (MOA; \citet{bon01}, \citet{sum03}) and the Optical Gravitational Lensing Experiment group (OGLE; \citet{uda03}). The MOA group uses the very wide field-of-view (2.2 square degrees) MOA-cam3 \citep{sak08} CCD camera mounted on the MOA-II 1.8m telescope at the Mt.\ John University Observatory in New Zealand. With this large FOV camera, MOA is able to observe 50 square degrees of our Galactic Bulge every hour, allowing high cadence observations. MOA detects about 600 new microlensing events and issues alerts of these events in real-time every year. The OGLE survey is conducted at the Las Campanas Observatory, Chile with the 1.3 m Warsaw telescope. In 2008, the OGLE was operating the OGLE-III survey using the 0.35 square degree OGLE-III camera, but OGLE has now upgraded to the 1.4 square degree OGLE-IV camera, which enables a higher cadence survey. This paper is a report of our analysis of a microlensing event OGLE-2008-BLG-355. The observations of this event are described in Section \ref{sec-obs}. Section \ref{sec-reduction} explains our data reduction procedure. Section \ref{sec-model} discusses our best model and the comparison with other models. The source color and the derived the source radius and the Einstein angular radius from the color are derivedf in Section \ref{sec-color}. The likelihood analysis is discussed in Section \ref{sec-like}. Finally, Section \ref{sec-disc} discusses the results of this work.	\label{sec-disc} This event was identified as a planetary event as a result of a systematic analysis of all past binary events observed by MOA prior to the 2013 bulge season. Previously, this event had been identified as a brown dwarf mass ratio binary lensing event using only OGLE data \citep{jar10}. Our systematic analysis finds that the best fit model using only the OGLE data is also a planetary model, and this points to the importance of a systematic analysis probing all of parameter space in order to find the best fit model. Our Bayesian likelihood analysis, based on a standard Galactic model indicates that the planet OGLE-2008-BLG-355Lb is a gas giant orbiting an M-dwarf or a late K-dwarf at 1$\sigma$ confidence. But the host could also be a G-dwarf. Such a massive planet with a mass ratio of $q = 0.0118 \pm 0.0006$, are predicted to be especially rare around low-mass stars, like M-dwarfs \citep{lau04,ken08}. Thus, one might be tempted to conclude that the existence of this planet is a challenge to the core accretion theory because an M-dwarf host star is favored by the Bayesian analysis. There is a flaw in this argument challenging the core accretion theory, however. Our Bayesian analysis assumed that host stars of all masses were equally likely to host a planet with the measured mass ratio, and so it could be that it is only this assumption that challenges the core accretion theory. To really test a theory, we need to start with a prior that is consistent with the theory, and then compare that prior to the data. A statistical analysis with planet detection efficiencies would be required to do a serious test of the theory. However, there has been no core accretion theory prediction of how the probability of hosting a planet of a given mass ratio should scale with the host star mass at the orbital separations probed by microlensing \citep{ida05}. The solution to this problem is to determine the host star mass. For some events \citep{gau08,ben08,mur11,kai13,pol13,tsa13,shv13}, this can be done with light curve measurements of finite source effects and the microlensing parallax effect, but the OGLE-2008-BLG-355 light curve does not allow a measurement of the microlensing parallax effect. Fortunately, lens star and planet masses can also be determined if the lens star is detected in high angular resolution follow-up observations \citep{ben06,bennett07,don09a,benet10,kub12,bat14}. In some cases, partial microlensing parallax information can be used to put constraints on the lens system mass \citep{bat11}, or a partial microlensing parallax measurement can be combined with high angular resolution follow-up observations \citep{don09b} to yield a lens system mass measurement. In the case of the two-planet system OGLE-2006-BLG-109Lb,c, the microlensing parallax mass measurement was confirmed by the host star detection in a high angular resolution image \citep{benet10}. Two planetary events similar to OGLE-2008-BLG-355 are OGLE-2003-BLG-235 and MOA-2011-BLG-293. In both cases, a planet with a super-Jupiter mass ratio ($q \gg 0.001$) was found orbiting a star determined to be a likely M-dwarf by a Bayesian analysis. In both cases, high angular resolution follow-up data was obtained after the event, and the follow-up data indicated that the lens stars were near the upper end of the mass range allowed by the Bayesian analysis. Neither host star turned out to be an M-dwarf. The host star OGLE-2003-BLG-235L was determined to have a mass of $M_h = 0.63 {+0.07\atop -0.09}\msun$ \citep{ben06}, and the host star MOA-2011-BLG-293L was found to have a mass of $M_h = 0.86\pm 0.06 \msun$ \citep{bat14}. This suggests that there may be some truth in the core accretion theory prediction that massive gas giants are rare around M-dwarfs, particularly low-mass M-dwarfs. The way to really test this core accretion theory prediction is to do a statistical analysis using events that have mass determinations from microlensing parallax measurements or high angular resolution follow-up observations that detect the host star. Table \ref{tab-events} lists the microlensing events with host mass determinations from either microlensing parallax or host star detection with high angular resolution follow-up observations. Figure~\ref{fig-like-color} shows that the host stars are likely to be within 3 magnitudes of the brightness of the source star in the $H$ or $K$-bands, based on the Bayesian analysis of lens system properties. But, if the core accretion theory prediction is right, then the lens star is likely to be on the bright side of the distributions in Figure~\ref{fig-like-color}, and so the lens star would be easier to detect than Figure~\ref{fig-like-color} implies. This event is also one that was characterized using only MOA and OGLE data, which were the survey groups active in 2008. There are several other planetary events which are characterized without any data from follow up groups \citep{bon04,ben08,yee12,ben12,pol13,shv13,suz14} and these planets are all gas giants except MOA-2007-BLG-192Lb, which has relatively sparse coverage over caustic but fortuitously can be characterized \citep{ben08}. Note that follow-up observations with NACO adaptive optics system on the VLT was conducted for MOA-2007-BLG-192 and the refined physical parameters of the lens system \citep{kub12} are consistent with the original results (see Table \ref{tab-events}). MOA's normal observation cadence for the field containing OGLE-2008-BLG-355 was every one observation per hour in 2008, but this event was characterized thanks to increases in cadence by both OGLE and MOA in response to the OGLE anomaly alert. At present, MOA has a 15 minute observing cadence in the 6 MOA fields ($13 {\rm deg}^2$) containing slightly more than half the microlensing events, while the OGLE-IV survey 3 fields ($4.2 {\rm deg}^2$) with a 20 minute cadence. These observing cadences should enable us to detect perturbations due to smaller planets such as cold Neptunes or even Earth-mass planets \citep{gau12}. Therefore, it is expected that the type of planetary systems may in the future be found more using only survey data. We acknowledge the following support: The MOA project was supported by a Grant-in-Aid for Scientific Research (JSPS19015005, JSPS19340058, JSPS20340052, JSPS20740104). D.P.B.\ was supported by grants NASA-NNX12AF54G and NSF AST-1211875. The OGLE project has received funding from the European Research Council under the European Community's Seventh Framework Programme (FP7/2007-2013)/ERC grant agreement no. 246678 to AU. %We acknowledge the following support: The Grant-in-Aid for Scientific Research...... (MOA project); The European Research Council under the European Community's Seventh Framework Programme (FP7/2007-2013)/ERC grant agreement no. 246678 to AU (OGLE project). %% The reference list follows the main body and any appendices. %% Use LaTeX's thebibliography environment to mark up your reference list. %% Note	14	3	1403.7005	We report the discovery of a massive planet, OGLE-2008-BLG-355Lb. The light curve analysis indicates a planet:host mass ratio of q = 0.0118 ± 0.0006 at a separation of 0.877 ± 0.010 Einstein radii. We do not measure a significant microlensing parallax signal and do not have high angular resolution images that could detect the planetary host star. Therefore, we do not have a direct measurement of the host star mass. A Bayesian analysis, assuming that all host stars have equal probability to host a planet with the measured mass ratio, implies a host star mass of M_{h} = 0.37_{-0.17}^{+0.30}\ M_{\odot } and a companion of mass M_{P} = 4.6^{+3.7}_{-2.2} M_{J}, at a projected separation of r_{\perp } = 1.70^{+0.29}_{-0.30} AU. The implied distance to the planetary system is D <SUB>L</SUB> = 6.8 ± 1.1 kpc. A planetary system with the properties preferred by the Bayesian analysis may be a challenge to the core accretion model of planet formation, as the core accretion model predicts that massive planets are far more likely to form around more massive host stars. This core accretion model prediction is not consistent with our Bayesian prior of an equal probability of host stars of all masses to host a planet with the measured mass ratio. Thus, if the core accretion model prediction is right, we should expect that follow-up high angular resolution observations will detect a host star with a mass in the upper part of the range allowed by the Bayesian analysis. That is, the host would probably be a K or G dwarf.	false	[ "host stars", "host mass ratio", "0.010 Einstein radii", "massive planets", "mass M_{P", "planet formation", "high angular resolution images", "a host star mass", "the host star mass", "more massive host stars", "the planetary host star", "0.010 Einstein", "equal probability", "a host star", "all host stars", "Bayesian", "the measured mass ratio", "follow-up high angular resolution observations", "1.70^{+0.29}_{-0.30", "This core accretion model prediction" ]	14.421645	12.374104	3
12366606	[ "Creminelli, Paolo", "Serone, Marco", "Trevisan, Gabriele", "Trincherini, Enrico" ]	2015JHEP...02..037C	[ "Inequivalence of coset constructions for spacetime symmetries" ]	29	[ "Abdus Salam International Centre for Theoretical Physics; , Institute for Advanced Study", "Abdus Salam International Centre for Theoretical Physics; , SISSA", "SISSA; , INFN — Sezione di Trieste", "Scuola Normale Superiore; , INFN — Sezione di Pisa" ]	[ "2014JHEP...10..006K", "2014JHEP...11..008D", "2014PhLB..733...46D", "2014PhRvD..90b4001B", "2014PhRvD..90b4050D", "2014PhRvD..90b5022G", "2014arXiv1403.6120B", "2015CQGra..32c5022D", "2015PhLB..744..347K", "2015PhRvD..92d4024M", "2015PhRvD..92d5020H", "2015arXiv150205706K", "2016PhRvD..93d5029B", "2016PhRvL.116d1601C", "2017JHEP...02..113C", "2017JHEP...02..134G", "2017JHEP...10..051K", "2017PhRvD..95f5019N", "2018EPJC...78..546M", "2018JHEP...03..051H", "2018JHEP...05..076B", "2019JHEP...11..166P", "2020JHEP...03..075F", "2020JHEP...12..056K", "2021JHEP...01..159G", "2022JHEP...04..128B", "2022PhRvD.106d3508S", "2024JHEP...06..004A", "2024arXiv240414518B" ]	[ "astronomy", "physics" ]	5	[ "Cosmology of Theories beyond the SM", "Conformal and W Symmetry", "SpaceTime Symmetries", "Sigma Models", "High Energy Physics - Theory", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1969PhRv..177.2239C", "1969PhRv..177.2247C", "1969PhRv..184.1760S", "1971AnPhy..62...98I", "1975TMP....25.1050I", "2000PhLB..485..208D", "2004JHEP...06..059N", "2006JHEP...10..014A", "2009PhRvD..79f4036N", "2010JCAP...05..015D", "2010JCAP...11..021C", "2010JHEP...05..095N", "2011JCAP...07..017G", "2012JCAP...12..030H", "2012JHEP...06..004G", "2012arXiv1212.6972C", "2013JHEP...10..040C", "2013JHEP...11..055N", "2013JMP....54f2503K", "2013PhRvL.110x1303H", "2014PhLB..733...46D", "2014PhRvD..89d5002N", "2014PhRvD..89f5006E", "2014PhRvD..89h5004B", "2014PhRvD..90b4050D" ]	[ "10.1007/JHEP02(2015)037", "10.48550/arXiv.1403.3095" ]	1403	1403.3095_arXiv.txt	The treatment of non-linear realizations of internal symmetries in quantum field theory is by now well understood. If a symmetry group $G$ is spontaneously broken to a subgroup $H$, the Goldstone theorem ensures that one massless spin zero particle occurs for each broken symmetry generator. Remarkably, the low-energy dynamics of Goldstone bosons is quite insensitive to the specific UV theory that originated them and is essentially governed by the group theory structure of the coset space $G/H$, as originally shown by Callan, Coleman, Wess and Zumino (CCWZ) \cite{Coleman:1969sm,Callan:1969sn}. For a given coset $G/H$, there is an infinite number of possible parameterizations leading to different Lagrangians, but these are all physically equivalent, since local field redefinitions do not change physical observables. The extension of such considerations to spacetime symmetries is far from trivial. First of all, the Goldstone theorem no longer applies and we lose the one-to-one correspondence between massless spin zero particles and broken generators. The number of Goldstone particles is less than the number of broken generators because certain constraints among the Goldstones can be imposed, the so-called ``inverse Higgs phenomenon" \cite{Ivanov:1975zq}.\footnote{ See \cite{Endlich:2013vfa,Brauner:2014aha} for recent analyses on the meaning of the inverse Higgs phenomenon.} A notable example is the spontaneous breaking of conformal symmetry down to the Poincar\'e group \cite{Salam:1970qk,Isham:1971dv}. In four dimensions the conformal group is isomorphic to SO(4,2) but of the naively expected 5 Goldstones only one is necessarily present, the dilaton. The technical generalization of the CCWZ coset construction to spacetime symmetries has been developed in \cite{Volkov:1973vd,Ivanov:1975zq}, but the physical interpretation of this construction is perhaps not totally understood. The spacetime coordinates enter now explicitly in the definition of the element of the coset. There is still in principle an infinite number of ways of parametrizing the relevant coset but finding explicit tractable parametrizations is not straightforward. Moreover, and most importantly for the present work, it is not clear whether and in what sense they are physically equivalent or not. Two relevant non-linear realizations of the conformal group (and their relation), dubbed DBI and Weyl representations in \cite{Creminelli:2013fxa}, have been derived in \cite{Bellucci:2002ji} from the coset construction of \cite{Volkov:1973vd,Ivanov:1975zq}. They are related by a complicated field and spacetime coordinate redefinition. We have recently shown that the relations between such representations can be seen as a change of coordinates in AdS$_5$ and lead to the same $2\rightarrow 2$ dilaton scattering amplitude \cite{Creminelli:2013fxa}. However, the complete physical equivalence of the two representations was not clear, since around non-trivial backgrounds the aperture of the lightcone of the dilaton in the two representations is different and theories which have only subluminal propagation are mapped into theories which allow superluminality. With this motivation, the aim of this paper is to shed light on the general question of whether different parametrizations of cosets involving spacetime symmetries are physically equivalent or not. We start in section \ref{sec:TwistAlgebra} by introducing what is going to be our prominent example of coset construction: Gal$(3+1,1)/$ISO$(3,1)$. This coset, introduced in \cite{Goon:2012dy} (see also \cite{Kamimura:2013mia}) to study Galileons \cite{Nicolis:2008in}, can be seen as the contraction of the conformal coset SO$(4,2)/$ISO$(3,1)$. It has the advantage, with respect to the latter, of being simpler to treat, yet keeping the same qualitative features. The two non-linear realizations of this coset that we consider are contractions of the DBI and Weyl representations of \cite{Bellucci:2002ji,Creminelli:2013fxa}. They are related to each other by a simple-looking field redefinition, eq.~(\ref{eq.coordtransf0}), that crucially involves the spacetime coordinates as well. Written as an ordinary field redefinition (i.e. keeping spacetime coordinates fixed), the field redefinition (\ref{eq.coordtransf0}) can be expanded in an infinite series of higher derivative terms. Recently the same field redefinition has also been found directly in the case of the standard (i.e. not conformal) Galileons in \cite{deRham:2013hsa} (see also \cite{Curtright:2012gx}). In fact for the coset Gal$(3+1,1)/$ISO$(3,1)$, since both the DBI and the Weyl operators reduce to the same set of five standard Galileons, the transformation that connects the two representations simply maps the five operators into themselves. In particular this gives rise to the apparent paradox presented in \cite{deRham:2013hsa}: a free massless scalar in one representation is mapped in the other into a non trivial Galileon theory in which superluminality seems to occur! In section \ref{sec.superluminal} we address this puzzle working at first order in the amplitude of the background. First of all we show that, in the theory obtained from the free theory by a field redefinition, there is no asymptotic effect if one considers a localized background which vanishes at infinity. The trajectory deviates from the lightcone, but the net effect integrates to zero asymptotically. This squares with the fact that the two theories have the same (trivial!) S-matrix, so that they should give the same answer to any question which involves asymptotic states. We then proceed to consider what happens if one measures the superluminality locally within a background, coupling the field to an external source, which does not transform when one changes the coset parametrization. The apparent paradox is solved by realizing that a local coupling to a source in one representation becomes non-local in the other representation. We analyze in detail how the non-local coupling provides a delay in the signal propagation that precisely compensates for the apparent superluminal behaviour. We thus conclude that the two representations have a different notion of locality and thus are not physically equivalent. In section \ref{sec.more} we provide several additional calculations to support this conclusion. In particular we show that there are no asymptotic effects at second order for a generic background in subsec.~\ref{sub.secondorder}. Furthermore we consider a cylindrically symmetric background in subsec. \ref{sub.cylinder}, and these two examples illustrate that the cancellation of asymptotic effects happens only for the particular linear combination of the five Galileons obtained by mapping the free theory. In subsec.~\ref{sub.sources} we show that our arguments apply even when sources are included. Then in subsec.~\ref{sub.DGP} we reconsider the Dvali-Gabadadze-Porrati (DGP) model \cite{Dvali:2000hr} and we verify we have asymptotic superlumnality in the case of a spherically symmetric solution. Finally, in subsec.~\ref{subsec:allorders}, we come back to our main example and show how luminality is recovered to all orders in the background amplitude for a specific, translationally invariant, classical configuration when non-local couplings are considered. In section \ref{sec:conformal} we show how the same mechanism applies to the conformal coset and explains the different aperture of the dilaton lightcone found in \cite{Creminelli:2013fxa} in the context of the Genesis scenario \cite{Creminelli:2010ba,Hinterbichler:2012yn}. In section \ref{sec:con} we summarize our findings and conclude with the general lesson that can be learned from these examples. For completeness, we report in appendix \ref{app:TdR} how the Galileon map found in \cite{deRham:2013hsa} is recovered through our coset construction. In appendix \ref{app:finite} we explain why the field redefinition we consider (a sum of an infinite number of higher derivative terms) cannot be truncated; if this is done either the superluminal effect is not measurable or we are forced to study higher derivative equations of motion.	\label{sec:con} In this paper we have studied two non-linear realizations of the Galileon group Gal($3+1$,1). These realizations are the contraction of the two known non-linear realizations of the conformal group, denoted DBI and Weyl representations in \cite{Creminelli:2013fxa}. The two representations are related by a field redefinition that crucially involves the spacetime coordinates as well. These field redefinitions (keeping the spacetime coordinates fixed) contain an infinite number of higher derivative terms. In the regime in which the series can be truncated to a certain number of terms, such as perturbative scattering processes involving a finite number of fields in a trivial background, nothing interesting happens: the two theories are simply equivalent. However, in discussing the propagation of perturbations around a background the series cannot be truncated. The propagation appears qualitatively different in the two representations. How is it possible for example that, starting from a free theory, one can induce propagation outside the Minkowski lightcone using a field redefinition? First of all, we showed that the two representations give the same result if one is interested in a localized background and in its asymptotic effect on a signal. This is not surprising since field redefinitions change the interactions but should not affect asymptotic properties, similarly to what happens to the S-matrix. However, if one is interested in local measurements in the presence of a background, the two representations still seem to give different answers about the superluminality with respect to the Minkowski lightcone. The point is that for this kind of measurements one must specify how fields couple with external sources and it turns out that local couplings in one representation become non-local in the other. Superluminality, starting from \cite{Adams:2006sv}, has been used as a way to rule out a conventional, local and Lorentz invariant UV completion. It is important to emphasize that an underlying assumption is that causality is defined by the Minkowski lightcone. This is typically taken for granted (and implicitly assumed in \cite{Adams:2006sv}), and is the reason why we insisted in keeping the external sources fixed (or equivalently the coordinates fixed). In this context, if one studies the Lagrangian for the field $q$ obtained from the free field theory without referring to the map, one can couple it in a generic (and local) way to other particles. In particular, one assumes that there can be other sectors of the theory which always propagate on the Minkowski lightcone, independently of any background. Clearly, this breaks the physical equivalence of the two representations: while in the initial free theory different massless particles move with the same speed, in the second theory different massless particles have different speed. With these assumptions, the $q$ theory is pathological since an asymptotic superluminality can be generated: even though the integrated effect for $q$ perturbations vanishes, one can obtain an overall effect using, in the regions where the $q$ speed is subluminal, other fields that propagate on the Minkowski lightcone. It is worth stressing an alternative procedure.\footnote{We thank C.~de Rham and A.~Tolley for discussions about this point and for sharing their work \cite{deRham:2014lqa} on the subject.} Instead of keeping the sources fixed, one could perform a redefinition of the source $J$ in eq.~(\ref{LJsource}): \be \int d^4y \, \pi(y) J(y) = \int d^4x\, {\rm Jac}(q(x))\, \Big(q(x) -\frac 1{2\Lambda^3} (\partial q)^2 \Big) \tilde J(x)\,, \ee where $\widetilde J(x) = J(y)$. In this way one does not induce any non-locality: we simply look at the change of coordinates (\ref{eq.coordtransf0}) as a diffeomorphism acting on all the fields of the theory. Interactions between all the fields and $q$ are induced in the mapped theory and as a consequence all the particles move on the same lightcone, the transformed of the Minkowski lightcone in the $\pi$ representation. In a general background for $q$, no particle is allowed to propagate on the Minkowski lightcone; they all propagate on an ``effective metric'' which depends on $q$ similarly to what happens in General Relativity and hence the Minkowski lightcone no longer defines causality. In this case it is meaningless to talk about local superluminality with respect to the Minkowski lightcone. We have simply rewritten the $\pi$ theory in a complicated way: in particular there is no asymptotic superluminality for any field. We conclude that, without further specifying the couplings with other fields and external sources, the presence of local superluminality does not imply pathologies, as shown by the example of the free theory mapped in the other representation. A less trivial example can be found by considering the conformal Galileons: starting from the Nambu-Goto action in the DBI representation (which is an interacting theory without pathologies), one can obtain the corresponding Weyl Lagrangian which admits local superluminal effects. On the other hand, asymptotic effects are absent in both representations: {\it only asymptotic superluminality is an unambiguous sign of pathology.} We concentrated on the propagation of perturbations, because this is a simple process where the relation between the two representations can be studied consistently without truncating the field redefinition. In this case the necessity of keeping an infinite number of operators, and thus be sensitive to non-locality, arises from having a large background field $\bar q$, which makes the condition of measurable superluminality, $\partial\bar q/\Lambda^2 \gtrsim 1$, compatible with the regime of validity of the EFT, $\partial/\Lambda \ll 1$. One may ask whether this has an S-matrix counterpart. In a perturbative computation of S-matrix elements with a given number of legs at any given order in perturbation theory, only a finite number of terms is relevant, and thus the series can be truncated. In this case the two representations give the same result and there is no sign of non-locality. On the other hand, we expect that the inequivalence of the two representations show up also in S-matrix elements in the absence of any background and where the series cannot be truncated, for example in the scattering of waves with large occupation number. It would be nice to have some concrete example of this. Our results about the inequivalence of different coset constructions for Gal($3+1$,1)/ISO(3,1) and SO(4,2)/ISO(3,1) should apply to other spacetime symmetries as well. It would be interesting to consider additional examples to provide further evidence to this claim, for example in the context of coset constructions at finite density \cite{Nicolis:2013sga} and for fluids \cite{Nicolis:2013lma}. As we have shown in our previous paper, the DBI and Weyl representations of the conformal group can both be obtained from a change of coordinates in AdS$_5$. From the AdS$_5$ point of view, roughly speaking, the first arises in the bulk (a brane in AdS$_5$) while the second arises on the boundary. It might be interesting to see if the non-local relation between the two representations can shed some light on some aspects of the AdS/CFT correspondence. \subsection*{Acknowledgements} We thank N. Arkani-Hamed, M.~Fasiello, R.~Penco, D.~Pirtskhalava, R.~Rattazzi, R.~Rosen, M.~Simonovi\'c and especially C.~de Rham, S.~Dubovsky, M.~Mirbabayi, A.~Nicolis, A.~Tolley and G.~Villadoro for useful discussions. P.C. acknowledges the support of the IBM Einstein Fellowship. The work of E.T. is supported in part by MIUR-FIRB grant RBFR12H1MW. \appendix	14	3	1403.3095	Non-linear realizations of spacetime symmetries can be obtained by a generalization of the coset construction valid for internal ones. The physical equivalence of different representations for spacetime symmetries is not obvious, since their relation involves not only a redefinition of the fields but also a field-dependent change of coordinates. A simple and relevant spacetime symmetry is obtained by the contraction of the 4D conformal group that leads to the Galileon group. We analyze two non-linear realizations of this group, focusing in particular on the propagation of signals around non-trivial backgrounds. The aperture of the lightcone is in general different in the two representations and in particular a free (luminal) massless scalar is mapped in a Galileon theory which admits superluminal propagation. We show that in this theory, if we consider backgrounds that vanish at infinity, there is no asymptotic effect: the displacement of the trajectory integrates to zero, as can be expected since the S-matrix is trivial. Regarding local measurements, we show that the puzzle is solved taking into account that a local coupling with fixed sources in one theory is mapped into a non-local coupling and we show that this effect compensates the different lightcone. Therefore the two theories have a different notion of locality. The same applies to the different non-linear realizations of the conformal group and we study the particular case of a cosmologically interesting background: the Galilean Genesis scenarios.	false	[ "non-trivial backgrounds", "Non-linear realizations", "backgrounds", "superluminal propagation", "different representations", "Galilean Genesis", "the different non-linear realizations", "internal ones", "spacetime symmetries", "two non-linear realizations", "a non-local coupling", "coordinates", "Galileon", "local measurements", "fixed sources", "signals", "infinity", "the 4D conformal group", "the Galileon group", "locality" ]	10.135906	0.417507	-1
525861	[ "Schwope, A. D.", "Thinius, B. D." ]	2014arXiv1403.4160S	[ "On the ephemeris of the eclipsing polar HU Aquarii" ]	0	[ "-", "-" ]	null	[ "astronomy" ]	6	[ "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1992ApJ...385..621A", "1993A&A...271L..25S", "1993MNRAS.263...61H", "1997A&A...319..894S", "2001A&A...375..419S", "2008AIPC..984..264V", "2009A&A...496..833S", "2011A&A...531A..34S", "2011MNRAS.414L..16Q", "2011MNRAS.416L..11H", "2012A&A...538A..84B", "2012AN....333..754P", "2012MNRAS.419.3258W", "2012MNRAS.420.3609H", "2012MNRAS.425..930G", "2013MNRAS.429L..45P" ]	[ "10.48550/arXiv.1403.4160" ]	1403	1403.4160_arXiv.txt	\hu\ is an eclipsing magnetic cataclysmic variable with a 125.0\,min orbital period. When discovered in 1993 as the optical counterpart to the soft X-ray and EUV sources RX\,J2107.9-0518/RE2107-05 \citep{hakala+93, schwope+93} it was the brightest eclipsing object displaying the most extended eclipse. Those properties triggered broad observational studies to disentangle accretion phenomena and the accretion geometry in a strongly magnetic environment. Particular emphasis was given to model the detailed eclipse structure. Comprehensive X-ray and EUV observations with the ROSAT and EUVE satellites took place between 1992 and 1998 \citep{schwope+01}. These studies established the eclipse egress as fiducial mark to determine the orbital period and a long-term ephemeris. The ingress into eclipse was found to be strongly affected by absorbing matter in the accretion curtain preventing an unequivocal determination of eclipse centre. The 31 epochs given for the eclipse egress that were based on soft X-ray observations already displayed systematic offsets with respect to a linear ephemeris of $\pm5$\,s. The size of the observed effect was still compatible with a migration of the accretion spot over the surface of the white dwarf. Monitoring observations with high-time resolution were continued at X-ray and optical wavelengths in the new millennium. \cite{schwarz+09} presented a further set of 62 eclipse epochs and were the first to discuss the timing residuals, that were then larger than the size of the white dwarf, in terms of an unseen third body and derived a possible mass of $M_3 \ga 5 M_{\rm Jup}$ for a planetary companion. \cite{qian+11} added a further 11 eclipse epochs and claimed the discovery of a circumbinary planetary system around the accreting binary. This hypothesis is debated. Their data seem to be offset from a more comprehensive data set presented by \cite{godz+12} and the proposed planetary system seems to loose stability on relatively short time scales \citep{horner+11,wittenmyer+12}. The study by \cite{hinse+12} on the other hand supports at least a two-planet scenario. Opportunities to reconstruct the evolution of the close binary through the presence of a planetary system were studied by \cite{port-zwart13}. The stellar and binary parameters were updated by \cite{schwope+11} who derive a total mass of the close binary of $0.98 \pm 0.10$\,\msun. The frequent low- and intermediate accretion states of the object prevented further detailed X-ray observations during the past 12 years. In July 2013 a high accretion state was noticed by monitoring observations with the robotic telescope STELLA at Tenerife. Swift and XMM-Newton observations were triggered to study the X-ray emission in a high accretion state. The observations with XMM-Newton were scheduled immediately after re-appearance of the source in the visibility zone of XMM-Newton and took place eventually on October 25. At that occasion attempts were made to obtain simultaneous ground-based data (Schwope et al.~2014, in preparation). Here we report quasi-simultaneous optical observations that were obtained utilizing the private equipment of the Inastars observatory Potsdam (IOP, B15, N 52.42392, E 13.012892).	The latest addition to the measured set of eclipse timings has revealed a very significant apparent decay of the orbital period. It is the first reported measurement of eclipse egress after two years without new data. We checked the reliability of the new data point by inspecting the data obtained quasi-simultaneously with XMM-Newton and found both data sets in agreement within a few seconds. Any possible mis-match of optical and X-ray timing data is too small to explain the huge trend in the $(O-C)$ data shown in both panels of Fig.~\ref{f:omc}. The results of the satellite data and a more thorough discussion of the relative timing between X-ray and optical data will be will be published elsewhere (Schwope et al.~2014, in preparation). The new data point lies outside the $(O-C)$ ranges of recent studies \citep{hinse+12,godz+12,wittenmyer+12}. \cite{hinse+12} were discussing a two-planet model with a summed amplitude of $23.4\pm0.1$\,s, too small, to make it compliant with our large $(O-C)$s. Their best-fit model also predicts increasingly later arrivals of the eclipse for cycle numbers $>$72000, contrary to what is observed. \cite{godz+12} were trying to fit the complete dataset with a two-planet model which gave apparently unphysical solutions only. Their linear model revealed massive third bodies, $M_3 \simeq 10$\,\msun. This model could probably be trimmed to reflect the large observed $(O-C)$ values, such a solution, however, would be mainly of academic interest. The quadratic model reveals an outer planet with eccentricity close to one, which was regarded unphysical by \cite{godz+12} themselves. A single planet model was found to be viable by ignoring data not obtained with the OPTIMA camera. This model with an amplitude of about 15\,s, however, is discarded by our result, which also discards the underlying assumption that the mix of data obtained at different wavelengths makes a meaningful fit so difficult. The planetary parameters reported in those studies therefore need revision. The current peak-to-peak amplitudes of the $(O-C)$ curves are 74\,s and 125\,s for the quadratic and linear fits, respectively. The LTT effect of a third body in a circular orbit seen edge-on is \begin{equation} \Delta T \simeq \frac{2M_3G^{1/3}}{c}\left(\frac{P_3}{2\pi(M_1+M_2)}\right)^{2/3}, \end{equation} with $M_1, M_2, M_3$ being the masses of the three bodies and $P_3$ the period of the third body around the binary \citep{pribulla+12}. Taken the measured peak-to-peak amplitudes at face value, planetary masses of 22 and 38 Jupiter masses for an assumed period of 6.9 years are derived. Such masses are significantly larger than the typically reported (non-pathological) planets around 5 Jupiter masses of past studies but should be treated with great caution; the parameters are not well constrained by our study. As the above equation shows, the mass scales with the timing amplitude which is not really measured but estimated from the current peak-to-peak amplitude of the timing residuals. The whole data set seems difficult to be reconciled with a single circumbinary planet due to the lack of regularity. On the other hand the data are not sufficient to finally constrain the parameters of a putative planetary system and to test the significance, the size, and physical interpretation of a quadratic term. This needs a much longer time base, in particular one needs to observe the turn-up towards later $(O-C)$ values. A note on the timing errors might be in place here. The current fit results are dominated by the OPTIMA data points with timing uncertainties as small as 0.11\,s. This is smaller than the accretion area, hence at that precision a timing jitter due to accretion filaments or other forms of in-stationary accretion becomes relevant. Without a correction from the observed time of eclipse egress to eclipse center or without an extra error term which accounts for the finite size of the accretion area, searches for planets might be biased by the high precision of some of the data, which do not reflect the LTT effect of a circumbinary planet (planetary system) but the effects of in-stationary accretion. The bottom line of this short communication is: The most recent addition to the $(O-C)$-data show an accelerated evolution of the timing residuals that rule out past models of a planet or a planetary system. HU Aqr deserves further monitoring with high cadence. \cite{godz+12} have ignored data obtained at other than optical wavelengths to make their one-planet model applicable. The measurement presented here shows that this self-restriction doesn't save a one-planet model. It appears more advantageous to use the complete data set for future modeling. The current data were obtained with private equipment operated in a semi-professional manner. They were obtained with a comparatively small telescope (by professional standards) located at a mildly light-polluted site. Nevertheless a highly constraining data point could be determined. More advanced amateurs with similar equipment might become interested to join the enterprise of searching for long-period planets via the LTT effect in collaboration with professionals. Potential targets and work in this direction are described by \cite{pribulla+12} and \cite{backhaus+12}. \acknowledgement We thank our referee for constructive criticism that help to improve the quality of the paper and Iris Traulsen and Robert Schwarz for help with the data reduction and useful discussions.	14	3	1403.4160	The magnetic cataclysmic variable HU Aquarii displayed pronounced quasi-periodic modulations of its eclipse timing. These were interpreted in terms of the light-travel time (LTT) effect caused by a circumbinary planet or planetary system. We report new photometric observations that revealed another precise eclipse timing for the October 2013 epoch, the first obtained in a high accretion state after many years in low or intermediate states. The eclipse was observed to occur earlier by 95.3 +- 2.0 s or 62.8 +- 2.0 s than expected for an assumed linear or quadratic ephemeris, respectively. The implied apparent strong evolution of the orbital period calls for a revision of the current planetary model or the planetary parameters. The object deserves further monitoring to uncover the true nature of the observed variability and to constrain the properties of the proposed planet or planetary system. The new observations prove that advanced amateur equipment can successfully be used in the growing field of planet search in wide circumbinary orbits via the LTT effect.	false	[ "many years", "low or intermediate states", "wide circumbinary orbits", "planet search", "a high accretion state", "pronounced quasi-periodic modulations", "a circumbinary planet or planetary system", "LTT", "the current planetary model", "the proposed planet or planetary system", "ephemeris", "the planetary parameters", "first", "advanced amateur equipment", "s", "the LTT effect", "HU Aquarii", "another precise eclipse timing", "new photometric observations", "further monitoring" ]	5.951908	10.554522	-1
483018	[ "Croll, Bryce", "Rappaport, Saul", "DeVore, John", "Gilliland, Ronald L.", "Crepp, Justin R.", "Howard, Andrew W.", "Star, Kimberly M.", "Chiang, Eugene", "Levine, Alan M.", "Jenkins, Jon M.", "Albert, Loic", "Bonomo, Aldo S.", "Fortney, Jonathan J.", "Isaacson, Howard" ]	2014ApJ...786..100C	[ "Multiwavelength Observations of the Candidate Disintegrating Sub-Mercury KIC 12557548b" ]	44	[ "Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA; NASA Sagan Fellow.;", "Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Visidyne, Inc., Santa Barbara, CA 93105, USA", "Center for Exoplanets and Habitable Worlds, The Pennsylvania State University, 525 Davey Lab, University Park, PA 16802, USA", "Department of Physics, University of Notre Dame, 225 Nieuwland Science Hall, Notre Dame, IN 46556, USA", "Institute for Astronomy, University of Hawaii at Manoa, 2680 Woodlawn Drive, Honolulu, HI 96822, USA", "Center for Exoplanets and Habitable Worlds, The Pennsylvania State University, 525 Davey Lab, University Park, PA 16802, USA", "Departments of Astronomy and of Earth and Planetary Science, University of California at Berkeley, Hearst Field Annex B-20, Berkeley, CA 94720-3411, USA", "Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "SETI Institute/NASA Ames Research Center, M/S 244-30, Moffett Field, CA 94035, USA", "Département de physique, Université de Montréal, C.P. 6128 Succ. Centre-Ville, Montréal, QC H3C 3J7, Canada", "INAF-Osservatorio Astrofisico di Torino, via Osservatorio 20, I-10025 Pino Torinese, Italy", "Department of Astronomy and Astrophysics, University of California, Santa Cruz, CA 95064, USA", "Department of Astronomy, University of California, Berkeley, CA 94720, USA" ]	[ "2014A&A...572A..76V", "2014AN....335.1018G", "2015AJ....149...24G", "2015AREPS..43..459T", "2015ApJ...800L..21B", "2015ApJ...802...28C", "2015ApJ...804...97C", "2015ApJ...809..139K", "2015ApJ...812..112S", "2015MNRAS.449.1408C", "2015MNRAS.454....2B", "2015Natur.526..546V", "2016A&A...589L...6A", "2016A&A...596A..32V", "2016ApJ...816...17W", "2016ApJ...826..156S", "2016ApJ...828...80K", "2016MNRAS.455.2018D", "2016MNRAS.461.1413V", "2016MNRAS.463.4422Z", "2017A&A...600A..86N", "2017AJ....154..242B", "2017ApJ...836...82C", "2017MNRAS.465.1008V", "2017MNRAS.469.3213H", "2017PASA...34...32K", "2018A&A...611A..63G", "2018A&A...618A..97R", "2018AJ....156..227C", "2018AJ....156..281S", "2018MNRAS.474.1453R", "2018MNRAS.474.4795X", "2018exha.book.....P", "2019A&A...628A..70R", "2019ApJ...876..127Z", "2019MNRAS.485.3876G", "2020rfma.book...45B", "2021AJ....162...57S", "2021arXiv210404824T", "2022ExA....53..729G", "2022NewAR..9401641L", "2022SSRv..218...15L", "2023MNRAS.518.4734S", "2024MNRAS.528.1249C" ]	[ "astronomy" ]	8	[ "eclipses", "infrared: planetary systems", "planetary systems", "stars: individual: KIC 12557548", "techniques: photometric", "Astrophysics - Earth and Planetary Astrophysics" ]	[ "1974SSRv...16..527H", "1977ApJ...217..425M", "1980ApJ...241L.105B", "1981lssp.book.....V", "1983asls.book.....B", "1988ApOpt..27.1203O", "1994SPIE.2198..362V", "1995A&A...300..503D", "1995Icar..114..203K", "1998RSEnv..66....1H", "2000PASP..112..315W", "2002Icar..159..529K", "2003A&A...408..193J", "2004ApJS..154...10F", "2004SPIE.5492..978P", "2006PASP..118.1351C", "2008ApJ...683.1076W", "2009Sci...325..709B", "2010A&A...521A..60M", "2010ApJ...711L..19B", "2010ApJ...717.1084C", "2010ApJ...718..920C", "2010ApJ...724.1108J", "2010MNRAS.408.1758K", "2011ApJ...729...27B", "2011JAtOT..28..779D", "2012A&A...537A.115B", "2012A&A...545L...5B", "2012ApJ...752....1R", "2012ApJ...761...39C", "2012arXiv1207.6212C", "2013A&A...557A..72B", "2013ApJ...776..134P", "2013ApJ...776L...6K", "2013Icar..222..634K", "2013MNRAS.433.2294P", "2013Natur.503..381H", "2014A&A...561A...3V", "2014ApJ...784...40R" ]	[ "10.1088/0004-637X/786/2/100", "10.48550/arXiv.1403.1879" ]	1403	1403.1879.txt	\citet{Rappaport12} presented intriguing, perplexing and downright peculiar {\it Kepler} observations of the K-dwarf star KIC 12557548. The {\it Kepler} space telescope's \citep{Borucki09} observations of this star displayed repeating dips every $\sim$15.7 hours that varied in depth from a maximum of $\sim$1.3\% of the stellar flux to a minimum of $\sim$0.2\% or less without a discernible rhyme or reason to explain the depth variations. In addition, the occultations were not the iconic transit-like shape we have come to expect from extrasolar planets or binary stars, but exhibited an obvious ingress/egress asymmetry, with a sharp ingress followed by a longer, more gradual egress. A non-detection of ellipsoidal light variations allowed an upper-limit on the mass of the occulting object to be set at 3 Jupiter masses, and thus prompted the question of what was causing this odd photometry? %After ruling out alternative possibilities of a planet precessing due to an unseen companion The answer the authors proposed was that the peculiar {\it Kepler} observations of KIC 12557548 (hereafter KIC 1255) are due to a gradually disintegrating low mass (super-Mercury) planet, KIC 12557548b (hereafter KIC 1255b). The thought process is that this putative planet, with its extremely short orbital period, is being roasted by its host star, and is throwing off material in fits and starts; at each passage in front of its parent star, the different amount of material being discarded by the planet leads to differences in the resulting optical depth, thus explaining the obvious transit depth variations. The clear ingress/egress asymmetry of the transit is then due to the fact that the material is streaming behind the planet, forming a long comet-like tail that obscures the star for a larger fraction of the orbit. %Something about the variable optical depth. Naturally, observations as odd as those presented by \citet{Rappaport12} and an explanation as exotic as a disintegrating super-Mercury, invited a great deal of skepticism from the astronomical community. Alternative theories that have been discussed to explain the observed photometry include: (i) a bizarre {\it Kepler} photometric artifact; (ii) a background blended eclipsing binary\footnote{Although how this would explain the ingress/egress asymmetry, or the transit depth variations remains a mystery.}; (iii) an exotically chaotic triple; or (iv) a binary that is orbiting KIC 12557548 wherein one member of the binary system is a white dwarf fed by an accretion disk \citep{Rappaport12}. Inspired by the \citet{Rappaport12} result there have been a number of modelling efforts to interpret the bizarre {\it Kepler} observations that seem to reinforce the possibility that the photometry of KIC 1255 is caused by scattering off material streaming from a disintegrating low-mass planet. Dust scattering models confirm the viability of the disintegrating planet scenario featuring a comet-like tail trailing the planet, composed of sub-micron sized grains \citep{Brogi12}, or up to one micron (0.1 - 1.0 $\mu m$) sized grains \citep{Budaj13}. These efforts suggest that the minute brightening just prior to transit % \footnote{We note we use the terms occultation and transit interchangeably in this draft.} can be readily explained by enhanced forward scattering from this dust cloud, while the ingress/egress asymmetry can be explained by a comet-like dust tail that has a particle density or size distribution that decreases with distance from the planet. The richness of the {\it Kepler} data on KIC 1255b have led to suggestions of evolution of the cometary tail \citep{Budaj13}, and that the comet is best explained by a two component model, with a dense coma and inner-tail and a diffuse outer tail \citep{Budaj13,vanWerkhoven13}. Another effort by \citet{Kawahara13} suggests that the observed transit depth variability may correlate with the stellar rotation period, and thus the presumed variable mass loss rate of the planet may be a byproduct of the stellar activity, specifically ultraviolet and X-ray radiation. \citet{PerezBeckerChiang13} argue from the results of a hydrodynamical wind model, that we may be observing the final death-throes of a planet catastrophically evaporating, and that KIC 1255b may range in mass from 0.02 - 0.07 $M_\oplus$ (less than twice that of the Moon, to greater than Mercury), although for most solutions the mass of KIC 1255b is less than that of Mercury. We note that subsequent to the submission of this work, there has been an announcement of a second low-mass planet candidate, possibly hosting a comet-like tail \citep{Rappaport14}. %INSERTHERE SHOULD I MENTION THAT I AM WORKING ON THIS. One proposed method for elucidating the unknown nature of the material that is supposedly occulting KIC 1255, is multiwavelength simultaneous observations of the transit of the object. As the efficiency of scattering diminishes for wavelengths longer than the approximate particle size \citep{HansenTravis74}, and given the inference of sub-micron size grains in the dust tail of this object \citep{Brogi12,Budaj13}, one might expect that infrared and near-infrared photometry of the transit of KIC 1255b would display significantly smaller depths than those displayed in the optical. Determining that the transit depth of KIC 1255b is wavelength dependent, with smaller depths in the near-infrared than the optical, would therefore strongly favour the explanation of scattering from a dust tail with sub-micron sized particles. Here we present an assortment of different observations of KIC 1255 that were obtained in order to either bolster or rule-out the disintegrating low-mass planet scenario. In addition to the {\it Kepler} photometry that we analyze here, these various observational data include: (i) two Canada-France-Hawaii Telescope/Wide-field InfraRed Camera (CFHT/WIRCam) Ks-band ($\sim$2.15 micron) photometric detections of the KIC 1255b transit with simultaneous {\it Kepler} photometric detections, (ii) simultaneous photometric non-detections of the KIC 1255b transit with the {\it Hubble Space Telescope} Wide Field Camera 3 ({\it HST}/WFC3) F140W and {\it Kepler} photometry, (iii) {\it HST}/WFC3 high-angular resolution imaging of KIC 1255 in the F555W ($\lambda$$\sim$0.531 $\mu m$), and F775W ($\lambda$$\sim$0.765 $\mu m$) bands, (iv) Keck/NIRC2 ground-based Adaptive Optics (tip/tilt only) high-angular resolution imaging of KIC 1255 in the K'-band ($\lambda\sim$2.124 $\mu m$), and (v) and Keck/HIRES radial velocity observations of KIC 1255. The high angular resolution imaging observations allow us to rule-out nearby background/foreground companions as close as 0.2\arcsec \ to KIC 1255. Our KECK/HIRES radial velocity observations allow us to rule-out low-mass stellar companions ($\sim$ 0.2 M$_\odot$) for orbital periods $\lesssim 10$ years. This significantly reduces the parameter space for nearby companions to KIC 1255, and therefore reduces the odds that the unique {\it Kepler} photometry that \citet{Rappaport12} reported is due to a binary or higher-order multiple, masquerading as a planetary false positive. Our simultaneous {\it Kepler} and near-infrared detections of the transit of KIC 1255b appear to report similar depths; as a result, if the source of the photometry we observe is a dust tail trailing a disintegrating planet composed of single-sized particles, then the particles are at least $\sim$\ParticleSizeCare \ microns in radius. Particles this large can likely only be lofted from a low-mass planet, suggesting that KIC 1255b might best be described as a sub-Mercury mass planet.	We have presented multiwavelength photometry, high angular-resolution imaging and radial velocities of the intriguing disintegrating low-mass candidate planet KIC 1255b. We summarize our findings here: (i) {\it Comparison of our CFHT/WIRCam 2.15 $\mu m$ to {\it Kepler} 0.6 $\mu m$ transit depths, and the resulting constraints on particle sizes in the tail trailing KIC 1255b:} The average ratio of the transit depths that we observe from the ground with CFHT/WIRCam and space with {\it Kepler} at our two epochs are: \RatioBOTHCFHTKepler \ $\pm$ \RatioBOTHErrorCFHTKepler. In the disintegrating planet scenario, the only way to see a lack of extinction from the optical to the near-infrared is if the circumference of the particles are at least approximately the wavelength of the observations. Therefore, if the transits we observe are due to scattering from single-size particles streaming from the planet in a comet-like tail, then the particles must be $\sim$\ParticleSizeCare \ microns in radius or larger\footnote{ We note this is in some disagreement, and modest agreement, with two efforts that presented scattering models compared to the {\it Kepler} photometry. The findings of \citet{Brogi12} modestly disagree with our own, as they suggest the particles in the tail trailing KIC 1255b must have a a typical grain size of 0.1 $\mu m$ from six quarters of long cadence {\it Kepler} photometry. The results of \citet{Budaj13} are in modest agreement with our own, as their analysis of 14 quarters of {\it Kepler} long and short cadence photometry, suggest grain sizes from 0.1 - 1.0 $\mu m$. }. Similarly, if the particle size distribution in the tail follows a number density defined by a power-law, then only smaller power-law slopes, and thus larger particle sizes result in a sufficient number of near-micron sized grains to satisfy our observations. (ii) {\it Comparison of our {\it HST} 1.4 $\mu m$ and {\it Kepler} 0.6 $\mu m$ null detections:} Unfortunately we were unable to detect the transit of KIC 1255b in either our simultaneous {\it HST} and {\it Kepler} photometry, due to the fact that the transits of KIC 1255b had largely disappeared in the {\it Kepler} photometry for approximately $\sim$5 days before and after our observed transit. We are therefore able to conclude little from these observations, other than there is no evidence for strongly different transit depths at these wavelengths. (iii) {\it Particle lifting from KIC 1255b:} \citet{PerezBeckerChiang13} have already demonstrated that lifting particles nearly a micron in size is possible from KIC 1255b. As lifting such large particles becomes much more difficult as one increases the mass of the candidate planet, we note our $\gtrsim$\ParticleSizeCare \ micron limit on single-sized particles in the tail trailing KIC 1255b favours a sub-Mercury, rather than super-Mercury, mass for KIC 1255b. (iv) {\it Constraints on false-positives from our high angular-resolution imaging, RVs and photometry:} Our {\it HST} ($\sim$0.53 $\mu m$ and $\sim$0.77 $\mu m$) high angular-resolution imaging allows us to rule-out background and foreground candidates at angular separations greater than 0.2\arcsec \ that could be responsible for the transit we associate with a planet transiting KIC 1255b. The associated limit from our groundbased Keck/NIRC2 Adaptive Optics observations in K'-band ($\sim$2.12 $\mu m$) is for separations greater than 1.4\arcsec. Our radial velocity observations allow us to rule out low-mass stellar companions ($\gtrsim$ 0.2 M$_\odot$) for periods less than $\lesssim$$10$ years, and 13 Jupiter-mass companions for periods less than $\lesssim$$40$ days. Furthermore, the similar transit depths we observe in the near-infrared with CFHT/WIRCam and in the optical with {\it Kepler} also allow us to rule out background/foreground candidates, or higher-order multiples with significantly different spectral types, as this would result in a colour-dependent transit depth from the optical to the near-infrared. Although prior to these observations we knew of no viable false-positive scenario that could reproduce the unique photometry we observed with {\it Kepler} (e.g. the forward scattering bump before transit, the sharp ingress and gradual egress transit profile, the sharply varying transit depths), we note that we have now greatly reduced the parameter space for viable false positive scenarios. We conclude that the disintegrating low-mass planet scenario is the simplest explanation for our multiwavelength photometry, RVs and high angular-resolution imaging suggested to date. (v) {\it Limit on the mass of the candidate planet KIC 1255b:} Our KECK/HIRES RVs of KIC 1255b allow us to place an upper-limit on the minimum mass of the candidate planet that confirms it is firmly in the planetary regime; this limit is $m \, \sin i$ $\lesssim$ 1.2 M$_J$ with 3$\sigma$ confidence, assuming a circular orbit. %PLANETMASSLIMIT	14	3	1403.1879	We present multiwavelength photometry, high angular resolution imaging, and radial velocities of the unique and confounding disintegrating low-mass planet candidate KIC 12557548b. Our high angular resolution imaging, which includes space-based Hubble Space Telescope Wide Field Camera 3 (HST/WFC3) observations in the optical (~0.53 μm and ~0.77 μm), and ground-based Keck/NIRC2 observations in K' band (~2.12 μm), allow us to rule out background and foreground candidates at angular separations greater than 0.''2 that are bright enough to be responsible for the transits we associate with KIC 12557548. Our radial velocity limit from Keck/HIRES allows us to rule out bound, low-mass stellar companions (~0.2 M <SUB>⊙</SUB>) to KIC 12557548 on orbits less than 10 yr, as well as placing an upper limit on the mass of the candidate planet of 1.2 Jupiter masses; therefore, the combination of our radial velocities, high angular resolution imaging, and photometry are able to rule out most false positive interpretations of the transits. Our precise multiwavelength photometry includes two simultaneous detections of the transit of KIC 12557548b using Canada-France-Hawaii Telescope/Wide-field InfraRed Camera (CFHT/WIRCam) at 2.15 μm and the Kepler space telescope at 0.6 μm, as well as simultaneous null-detections of the transit by Kepler and HST/WFC3 at 1.4 μm. Our simultaneous HST/WFC3 and Kepler null-detections provide no evidence for radically different transit depths at these wavelengths. Our simultaneous CFHT/WIRCam detections in the near-infrared and with Kepler in the optical reveal very similar transit depths (the average ratio of the transit depths at ~2.15 μm compared with ~0.6 μm is: 1.02 ± 0.20). This suggests that if the transits we observe are due to scattering from single-size particles streaming from the planet in a comet-like tail, then the particles must be ~0.5 μm in radius or larger, which would favor that KIC 12557548b is a sub-Mercury rather than super-Mercury mass planet. <P />Based on observations obtained with WIRCam, a joint project of CFHT, Taiwan, Korea, Canada, and France, at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institute National des Sciences de l'Univers of the Centre National de la Recherche Scientifique of France, and the University of Hawaii. <P />Some of the data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. The observatory was made possible by the generous financial support of the W.M. Keck Foundation.	false	[ "super-Mercury mass planet", "KIC 12557548b", "KIC", "low-mass planet candidate KIC", "Hubble Space Telescope Wide", "Canada", "France", "high angular resolution imaging", "μm", "radial velocities", "InfraRed Camera", "Hawaii", "CFHT", "space-based Hubble Space Telescope Wide Field Camera", "Kepler", "angular separations", "~2.15 μm", "radically different transit depths", "very similar transit depths", "the Centre National de la Recherche Scientifique" ]	6.875651	13.918154	-1
678488	[ "Feldstein, Brian", "Kahlhoefer, Felix" ]	2014JCAP...08..065F	[ "A new halo-independent approach to dark matter direct detection analysis" ]	62	[ "Rudolf Peierls Centre for Theoretical Physics, University of Oxford, 1 Keble Road, Oxford OX1 3NP, United Kingdom", "Rudolf Peierls Centre for Theoretical Physics, University of Oxford, 1 Keble Road, Oxford OX1 3NP, United Kingdom" ]	[ "2014JCAP...09..049C", "2014JCAP...10..022C", "2014JCAP...10..076F", "2014JCAP...12..015B", "2014JCAP...12..052F", "2014arXiv1411.4557K", "2015JCAP...03..012D", "2015JCAP...07..019K", "2015JCAP...08..039B", "2015JCAP...08..046D", "2015JCAP...09..012H", "2015JCAP...09..052F", "2015JCAP...10..012A", "2015JCAP...11..038G", "2015PhPro..61...45D", "2015PhRvD..91a5017B", "2015PhRvD..91j3533K", "2015PhRvL.114w1303C", "2016JCAP...03..019M", "2016JCAP...05..024B", "2016JCAP...10..029G", "2016JCAP...10..032K", "2016JPhCS.718d2063W", "2016PhRvD..93e5009H", "2016PhRvD..94f3527O", "2016PhRvD..94l3009K", "2017JCAP...05..026W", "2017JCAP...08..039I", "2017JCAP...09..032G", "2017JCAP...10..002F", "2017JCAP...11..016K", "2017JCAP...12..039G", "2017JHEP...10..059K", "2017JPhG...44h4001G", "2017PhRvD..95f3017O", "2018IJMPA..3350120B", "2018JCAP...04..052H", "2018JCAP...05..074K", "2018JCAP...07..028C", "2018JCAP...12..018I", "2018PhLB..779..388K", "2018PhLB..785..610M", "2018PhRvD..98j3006O", "2018arXiv180102905K", "2018arXiv181011468E", "2019JCAP...12..013I", "2019PDU....2600393H", "2019PhRvD..99b3012E", "2019PhRvD..99d3541E", "2019arXiv190607541E", "2020PhRvD.101b3006O", "2020arXiv200812587V", "2021JCAP...12..048C", "2022LNP...996.....D", "2023JCAP...03..011K", "2023JCAP...11..077K", "2023JHEP...02..200M", "2023PhLB..84137922C", "2023arXiv231001480L", "2023arXiv231001483L", "2024JCAP...03..047B", "2024arXiv240304959H" ]	[ "astronomy", "physics" ]	3	[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1990NIMPA.297..496B", "1996APh.....6...87L", "1998PhRvD..57.3256K", "2001PhRvD..64d3502S", "2007JCAP...06..011D", "2008JCAP...06..012D", "2008Natur.454..735D", "2009JCAP...04..010S", "2009JCAP...11..019S", "2009MNRAS.397...44R", "2010EPJC...67...39B", "2010JCAP...01..006C", "2010JCAP...01..020F", "2010JCAP...02..030K", "2010JCAP...10..034G", "2010PhRvD..81h7301P", "2010PhRvD..82b3530M", "2010PhRvD..82f3512G", "2010arXiv1004.0691E", "2011JCAP...09..022A", "2011JCAP...10..035P", "2011PhRvD..83b3519L", "2011PhRvD..83c4007F", "2011PhRvD..83h3505P", "2011PhRvD..83j3514F", "2011PhRvD..83l5029P", "2012JCAP...01..024F", "2012JCAP...08..010D", "2012JCAP...12..015G", "2012PhRvD..86f5027K", "2013APh....47...45F", "2013JCAP...02..041P", "2013JCAP...10..048D", "2013PhRvL.111c1302K", "2013PhRvL.111y1301A", "2014PDU.....5...45P", "2014PhRvD..89h5026K", "2014PhRvL.112i1303A", "2014PhRvL.112j1301F", "2014PhRvL.112x1302A", "2014arXiv1401.3295A" ]	[ "10.1088/1475-7516/2014/08/065", "10.48550/arXiv.1403.4606" ]	1403	1403.4606_arXiv.txt	The direct search for dark matter (DM) in shielded underground detectors is a promising strategy not only for confirming the particle nature of DM, but also for measuring its key properties, such as mass and couplings to nucleons. A central problem in the analysis of such experiments, however, is the uncertainty in the DM velocity distribution $f(\mathbf{v})$~\cite{Kamionkowski:1997xg,Green:2010gw, McCabe:2010zh}. Numerical simulations indicate that assuming an isotropic and isothermal halo may not be a good approximation~\cite{Kuhlen:2009vh} and that in addition there may be localised streams of DM~\cite{Diemand:2008in} as well as a DM disk co-rotating with the stars~\cite{Read:2009iv}. Such uncertainties are particularly important for light DM, which only probes the tail of $f(\mathbf{v})$~\cite{Lisanti:2010qx, Fairbairn:2012zs}, but they also significantly reduce the amount of information that can be inferred about general DM candidates. In fact, since changes in the halo structure can mimic changes in the DM parameters, a single direct detection experiment is insufficient to determine even the DM mass \cite{Drees:2008bv}. To make progress, one has to find a way to combine information from different target materials and quantify the impact of astrophysical uncertainties~\cite{Pato:2010zk, Pato:2011de, Peter:2013aha}. The standard approach to this problem is to parametrize the uncertainties in $f(\mathbf{v})$ and scan (or marginalise) over the associated parameters~\cite{Strigari:2009zb, Peter:2009ak, Peter:2011eu, Arina:2011si, Pato:2012fw, Kavanagh:2012nr, Kavanagh:2013wba} (for alternative strategies see~\cite{Drees:2007hr, Drees:2008bv, Fox:2010bz}). This approach suffers from the problem that it is unclear whether the chosen parameterisation of the halo is sufficiently general. Moreover, direct detection experiments do not probe $f(\mathbf{v})$ directly, but instead probe the velocity integral $g(v_\text{min}) = \int_{v>v_\text{min}} \hspace{-1mm} f(\mathbf{v}) / v \, \mathrm{d}^3\mathbf{v} $. Therefore, it is typically necessary to perform large numbers of numerical integrations over the DM velocity distribution, making this method numerically slow. Efficient scans then often require a Bayesian approach, with the need to motivate prior distributions for DM parameters. In this letter we propose a new method for dealing with uncertainties in the DM velocity distribution. Instead of parametrizing $f(\mathbf{v})$, we directly parametrize $g(v_\text{min})$, so that predicted event rates depend on the parameters in a very simple way. In analogy to the treatment of the DM speed distribution $f(v) = \int f(\mathbf{v}) \mathrm{d}\Omega_v$ in~\cite{Peter:2009ak, Peter:2011eu}, we will write the velocity integral as a sum of step functions. This approach allows us to have a very large number of free parameters and removes the need to make any assumptions about the form of $f(\mathbf{v})$. Consequently, any conclusions drawn from it will be completely robust in the face of astrophysical uncertainties. Moreover, our method involves a frequentist rather than a Bayesian approach, so that no prior distributions for any of the parameters need to be proposed.		14	3	1403.4606	Uncertainty in the local dark matter velocity distribution is a key difficulty in the analysis of data from direct detection experiments. Here we propose a new approach for dealing with this uncertainty, which does not involve any assumptions about the structure of the dark matter halo. Given a dark matter model, our method yields the velocity distribution which best describes a set of direct detection data as a finite sum of streams with optimised speeds and densities. The method is conceptually simple and numerically very efficient. We give an explicit example in which the method is applied to determining the ratio of proton to neutron couplings of dark matter from a hypothetical set of future data.	false	[ "direct detection data", "direct detection experiments", "future data", "dark matter", "data", "optimised speeds", "densities", "the local dark matter velocity distribution", "neutron couplings", "streams", "the dark matter halo", "a dark matter model", "proton", "a finite sum", "Uncertainty", "the velocity distribution", "a hypothetical set", "a key difficulty", "a set", "The method" ]	7.767682	-2.229542	54
578138	[ "Stoyanov, K. A.", "Zamanov, R. K.", "Bode, M. F.", "Pritchard, J.", "Tomov, N. A.", "Gomboc, A.", "Koleva, K." ]	2014BlgAJ..21...32S	[ "Emission line variability in the spectrum of V417 Centauri" ]	1	[ "Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of Sciences, 72 Tsarighradsko Shousse Blvd., 1784 Sofia, Bulgaria", "Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of Sciences, 72 Tsarighradsko Shousse Blvd., 1784 Sofia, Bulgaria", "Astrophysics Research Institute, Liverpool John Moores University, IC2 Liverpool Science Park, 146 Brownlow Hill, Liverpool L3 5RF, UK", "European Southern Observatory, Karl-Schwarzschild-Str. 2, D-85748 Garching bei München, Germany", "Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of Sciences, 72 Tsarighradsko Shousse Blvd., 1784 Sofia, Bulgaria", "Faculty of Mathematics and Physics, University of Ljubljana, Jadranska 19, 1000 Ljubljana, Slovenia and Centre of Excellence SPACE-SI, A\\vskerčeva cesta 12, SI-1000, Ljubljana, Slovenia", "Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of Sciences, 72 Tsarighradsko Shousse Blvd., 1784 Sofia, Bulgaria" ]	[ "2021BlgAJ..34....3Z" ]	[ "astronomy" ]	4	[ "stars: individual: V417 Cen", "binaries: symbiotic", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1977A&A....57..383Z", "1981A&A....99..126H", "1986AJ.....92.1118K", "1988AJ.....95...45A", "1988ASPC....1...67S", "1989A&A...220..112Z", "1992A&A...265L..37S", "1994A&A...285..241V", "1994A&AS..106..243C", "1996ApJ...470.1187R", "1997ppsb.conf..147C", "1999A&A...347..151T", "1999Msngr..95....8K", "1999anot.conf..192C", "2002A&A...383..188M", "2002MNRAS.329..897H", "2002MNRAS.337.1038S", "2003ASPC..303.....C", "2004PASJ...56..353I", "2005A&A...441.1135J", "2005ARA&A..43..435H", "2006MNRAS.365.1215Z", "2007A&A...467.1389C", "2007A&A...472..497A", "2007BaltA..16..104F", "2010IAUS..262..307B", "2012MNRAS.423L..35P", "2013SCPMA..56..663A" ]	[ "10.48550/arXiv.1403.6603" ]	1403	1403.6603_arXiv.txt	Symbiotic stars (SSs) are thought to comprise a compact object (usually a white dwarf) accreting from a cool giant. They offer a laboratory in which to study such processes as (1) mass loss from cool giants and the formation of nebulae, (2) accretion onto compact objects, radiative transfer in gaseous nebulae, (3) jets and outflows (i.e. Corradi, Mikolajewska \& Mahoney 2003).\\ V417~Centauri (HV~6516, Hen~3-977) is a poorly studied D'-type (yellow) symbiotic system surrounded by a faint asymmetric nebula. The symbiotic nature of the object was proposed by Steiner, Cieslinski \& Jablonski (1988). The cool component is a G2~Ib-II star, with $ \log (L/L_\odot) =3.5$ and $T_{eff}=5000$~K (van Winckel et al. 1994). Pereira, Cunha \& Smith (2003) found atmospheric parameters $T_{eff} = 6000$, $\log g =3.0$, and spectral type F9 III/IV. The binary period is not defined. Van Winckel et al. (1994) found a 245.68~day periodicity with an amplitude of 0.5 mag using Harvard and Sonneberg plates. Gromadzki et al. (2011) using optical photometric observations covering 20 years detected strong long term modulation with a period of about 1700 days and amplitude about 1.5 mag in V-band, in addition to variations with shorter time-scales and lower amplitudes. However, the long period seems to be non-coherent and the nature of light variations and the length of the orbital period remain unknown. Here we discuss the emission line variability of V417~Cen in 2004.	The Balmer emission lines, appearing in the extended envelopes of symbiotic stars, usually have an ordinary nebular profile with typical FWHM~$\approx$~100-150~\kms, for example AG~Dra (Tomova \& Tomov 1999). FWHM increases to 200~\kms\ only during the active phase. The basic mechanism determining their width is turbulence in the gas. The appearance and the variability of the narrow component is relatively common in symbiotic stars (see also Ikeda \& Tamura 2004 and the references therein). The central narrow emission of V417~Cen is very similar. When the star brightness increases, the EW of the line decreases. The FWHM and line fluxes of the Balmer lines remain unchanged. Our interpretation is that the continuum of G2~Ib-II supergiant was obscured by dust and after it the dust absorption decreased. In other words, during the time of the light minimum the G9 supergiant spectrum is veiled by dust. Rapid rotation is a common property of the cool components of D' SSs (Jorissen et al. 2005, Zamanov et al. 2006). This fast rotation is due to the mass transfer (spin accretion from the former AGB) and/or the tidal force (Soker 2002; Ablimit \& L\"u 2012). Fast rotation in D'-type symbiotics probably leads to the formation of a circumstellar disk as in the classical Be stars, where the fast rotation of the B star expels matter in the equatorial regions (see Porter \& Rivinius 2003). The geometry and kinematics of the circumstellar environment in Be stars is best explained by a rotationally supported relatively thin disk with very little outflow. The central B star is a fast rotator with a commonly quoted mean value of about 70\% - 80\% of the critical velocity, i.e. the ratio $v_{rot}/v_{cr} \sim 0.7$. Regarding the ratio $v_{rot}/v_{cr} \sim 0.7$ the rotation of the G2Ib-II component is similar to the Be stars. This is estimated as $ v \, sin \, i = 75$ \kms, which is 71\% of the critical value and implies a short rotation period of the mass donor $P_{rot} \le 51$~d (Zamanov et al. 2006). By analogy with the Be stars and Be/X-ray binaries, we expect a rotationally supported disk around the supergiant in V417~Cen with little outflow. The far-IR data (Kenyon, Fernandez-Castro \& Stencel 1988) for D'-type symbiotics indicate optically thick dust clouds, with temperature $T_{dust} \sim 100 - 400$~K (Kenyon, Fernandez-Castro \& Stencel 1988). The IR colours of V417~Cen reveal the presence of warmer dust more similar to D-type symbiotics. The IR energy distribution peaks in the L band (3.79 $\mu m$), corresponding to a dust temperature of 800 ~K. The infrared excess is broad and cannot be fitted with a single black body, indicating temperature stratification in the dust (van Winckel et al. 1994). Angeloni et al. (2007), for another D'-type symbiotic HD330036, detected three dust shells, which are probably circumbinary. The IR energy distribution of V417~Cen is similar to the distribution in D-type symbiotics with resolved bipolar nebulae like BI~Cru, He 2-104 and R~Aqr (Schwarz \& Corradi 1992), with a broad IR excess peaking in the L band. In these systems, an excretion disk is thought to provide equatorial density enhancement. Where is the dust located? There are three suggestions for this: \begin{itemize} \item The hot component lies outside the dust shell that enshrouds the giant companion (as suggested for D-type symbiotics by Kenyon, Fernandez-Castro \& Stencel 1986). \item The dust is located around the hot component and at $L_4$ and $L_5$ Lagrangian points (Gromadzki et al. 2011). \item Dust shells are circumbinary, i.e. the orbital period is short and the supergiant and WD are in the centre of the dust shells, as supposed by Angeloni et al. (2007) for HD330036. \end{itemize}	14	3	1403.6603	We report high resolution (λ/Δ λ ∼ 48000) spectral observations of the yellow symbiotic star V417 Cen obtained in 2004. We find that the equivalent widths of the emission lines decreased, while the brightness increased. The FWHMs and wavelengths of the emission lines do not change. <P />We estimated the interstellar extinction towards V417 Cen as E_{B-V} = 0.95 ±0.10, using the KI interstellar line. <P />Using the [O III] lines, we obtain a rough estimation of the density and the temperature in the forbidden lines region log N_e ≈ 4.5 ± 0.5 and T_e=100000 ± 25000 K. Tidal interaction in this binary is also discussed.	false	[ "V417 Cen", "N_e ≈", "the KI interstellar line", "the forbidden lines region", "KI", "T_e=100000", "spectral observations", "the emission lines", "V417 Cen as E_{B-V", "the yellow symbiotic star", "the interstellar extinction", "high resolution", "a rough estimation", "Δ", "4.5 ±", "E_{B-V} = 0.95 ±0.10", "λ", "this binary", "wavelengths", "the density" ]	6.977571	10.602884	-1
524430	[ "Fan, R. R.", "Zhang, F.", "Peng, W. X.", "Dong, Y. F.", "Gong, K.", "Yang, S.", "Guo, D. Y.", "Wang, J. Z.", "Gao, M.", "Liang, X. H.", "Zhang, J. Y.", "Cui, X. Z.", "Liu, Y. Q.", "Wang, H. Y." ]	2014arXiv1403.1679F	[ "The silicon matrix for the prototype for the Dark Matter Particle Explorer" ]	0	[ "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-" ]	null	[ "astronomy", "physics" ]	5	[ "Physics - Instrumentation and Detectors", "Astrophysics - Instrumentation and Methods for Astrophysics", "High Energy Physics - Experiment" ]	[ "2008Natur.456..362C", "2009Natur.458..607A", "2012PhRvL.108a1103A", "2013PhRvL.110n1102A" ]	[ "10.48550/arXiv.1403.1679" ]	1403	1403.1679_arXiv.txt	\label{Introduction} With the results of AMS-02\cite{1}, ATIC\cite{2}, PAMELA\cite{3} and FERMI\cite{4}, more and more proofs pointed the electron excess of cosmic electrons. With a requirement of higher and more precise energy measurement, a dark matter search satellite project DAMPE has been put forward. The main scientific objective of DAMPE is to measure electrons and photons with much higher energy resolution and energy reach than achievable with existing space experiments in order to identify possible Dark Matter signatures. It has also great potential in advancing the understanding of the origin and propagation mechanism of high energy cosmic rays, as well as in new discoveries in high energy gamma astronomy. There are four sub-detector in the DAMPE system, the plastic scintillator(PSD), the silicon tracker(STK), the BGO electromagnetic calorimeter(BGO) and the neutron detector.(see Figure \ref{DAMPE.fig}) \begin{figure} \includegraphics[width=5in]{DAMPE.png} \caption{The schematic of DAMPE detector: The components of the DAMPE are PSD, STK, BGO and the neutron detector from top to bottom.}\label{DAMPE.fig} \end{figure} The plastic scintillator and silicon tracker are employed to charge particles' discrimination and reconstruction the track. The silicon tracker with tungsten convertors can also distinguish the gamma ray and electrons. The STK is followed by an imaging calorimeter of about 31 radiation lengths thickness, made up of 14 layers of BGO bars in a hodoscopic arrangement. A layer of neutron detectors is added to the bottom of the calorimeter for the proton/electron discrimination. The total thickness of the BGO calorimeter and the STK correspond to about 33 radiation lengths, making it the deepest calorimeter ever used in space. For electrons and photons, the detection range of the DAMPE is 5 GeV - 10 TeV, with an energy resolution of about 1\% at 800 GeV. For cosmic rays, the detection range is 100 GeV - 100 TeV, with an energy resolution better than 40\% at 800 GeV. The geometrical factor is about ${0.3 m^{2}sr}$ for electrons and photons, and about ${0.2m^{2}sr}$ for cosmic rays. The angular resolution is 0.1$^\circ$ at 100 GeV. During the DAMPE prototype testing, a silicon matrix built by the Institute of High Energy Phsics(IHEP), the Chinese Academy of Sciences(CAS), serves as the silicon tracker. Followed chapters describe the details of the silicon matrix construction and test.	\label{Conclusion} With the cosmic ray test, the matrix can distinguish the MIPs signal from the electronics noise clearly, while the position can be confirmed by the channel sequence. Through a long time running, the matrix is optimized for the DAMPE prototype. The results of the DAMPE prototype test will be presented in near future.	14	3	1403.1679	A new generation detector for the high energy cosmic ray - the DAMPE(DArk Matter Particle Explorer) is a satellite based project. Its main object is the measurement of energy spectrum of cosmic ray nuclei from 100GeV to 100TeV, the high energy electrons and gamma ray from 5GeV to 10TeV. A silicon matrix detector described in this paper, is employed for the sea level cosmic ray energy and position detection while the prototype testing of the DAMPE. This matrix is composed by the 180 silicon PIN detectors, which covers an area of 32*20 cm2. The primary testing results are shown including MIPs energy spectrum and the position sensitive map.	false	[ "cosmic ray", "cosmic ray nuclei", "MIPs energy spectrum", "energy spectrum", "gamma ray", "Particle Explorer", "the sea level cosmic ray energy and position detection", "the high energy electrons", "a satellite based project", "the high energy", "the position sensitive map", "MIPs", "The primary testing results", "PIN", "A silicon matrix detector", "A new generation detector", "the DAMPE", "the 180 silicon PIN detectors", "Its main object", "an area" ]	7.255002	-2.000201	-1
745745	[ "Brothwell, R. D.", "Watson, C. A.", "Hébrard, G.", "Triaud, A. H. M. J.", "Cegla, H. M.", "Santerne, A.", "Hébrard, E.", "Anderson, D. R.", "Pollacco, D.", "Simpson, E. K.", "Bouchy, F.", "Brown, D. J. A.", "Chew, Y. Gómez Maqueo", "Collier Cameron, A.", "Armstrong, D. J.", "Barros, S. C. C.", "Bento, J.", "Bochinski, J.", "Burwitz, V.", "Busuttil, R.", "Delrez, L.", "Doyle, A. P.", "Faedi, F.", "Fumel, A.", "Gillon, M.", "Haswell, C. A.", "Hellier, C.", "Jehin, E.", "Kolb, U.", "Lendl, M.", "Liebig, C.", "Maxted, P. F. L.", "McCormac, J.", "Miller, G. R. M.", "Norton, A. J.", "Pepe, F.", "Queloz, D.", "Rodríguez, J.", "Ségransan, D.", "Skillen, I.", "Smalley, B.", "Stassun, K. G.", "Udry, S.", "West, R. G.", "Wheatley, P. J." ]	2014MNRAS.440.3392B	[ "A window on exoplanet dynamical histories: Rossiter-McLaughlin observations of WASP-13b and WASP-32b" ]	19	[ "Astrophysics Research Centre, School of Mathematics & Physics, Queen's University, Belfast BT7 1NN, UK", "Astrophysics Research Centre, School of Mathematics & Physics, Queen's University, Belfast BT7 1NN, UK", "Institut d'Astrophysique de Paris, UMR7095 CNRS, Université Pierre & Marie Curie, 98bis boulevard Arago, F-75014 Paris, France; Observatoire de Haute-Provence, CNRS/OAMP, F-04870 St Michel l'Observatoire, France", "Observatoire astronomique de l'Université de Genève, 51 ch. des Maillettes, CH-1290 Sauverny, Switzerland; Department of Physics, and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Astrophysics Research Centre, School of Mathematics & Physics, Queen's University, Belfast BT7 1NN, UK", "Centro de Astrofisica, Universidade do Porto, Rua das Estrellas, P-4150-762 Porto, Portugal", "IRAP-UMR 5277, CNRS and Univ. de Toulouse, 14 Av. E. Belin, F-31400 Toulouse, France", "Astrophysics Group, Keele University, Staffordshire ST5 5BG, UK", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK", "Astrophysics Research Centre, School of Mathematics & Physics, Queen's University, Belfast BT7 1NN, UK", "Institut d'Astrophysique de Paris, UMR7095 CNRS, Université Pierre & Marie Curie, 98bis boulevard Arago, F-75014 Paris, France; Observatoire de Haute-Provence, CNRS/OAMP, F-04870 St Michel l'Observatoire, France", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK", "School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews, Fife KY16 9SS, UK", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, F-13388 Marseille, France", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK; Department of Physics and Astronomy, Macquarie University, NSW 2109, Australia", "Department of Physical Sciences, The Open University, Milton Keynes MK7 6AA, UK", "Max Planck Institut für Extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany", "Department of Physical Sciences, The Open University, Milton Keynes MK7 6AA, UK", "Institut d'Astrophysique et Géophysique, Université de Liège, allée du 6 Août 17, B-4000 Liège, Belgium", "Astrophysics Group, Keele University, Staffordshire ST5 5BG, UK", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK", "Institut d'Astrophysique et Géophysique, Université de Liège, allée du 6 Août 17, B-4000 Liège, Belgium", "Institut d'Astrophysique et Géophysique, Université de Liège, allée du 6 Août 17, B-4000 Liège, Belgium", "Department of Physical Sciences, The Open University, Milton Keynes MK7 6AA, UK", "Astrophysics Group, Keele University, Staffordshire ST5 5BG, UK", "Institut d'Astrophysique et Géophysique, Université de Liège, allée du 6 Août 17, B-4000 Liège, Belgium", "Department of Physical Sciences, The Open University, Milton Keynes MK7 6AA, UK", "Observatoire astronomique de l'Université de Genève, 51 ch. des Maillettes, CH-1290 Sauverny, Switzerland", "School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews, Fife KY16 9SS, UK", "Astrophysics Group, Keele University, Staffordshire ST5 5BG, UK", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK; Isaac Newton Group of Telescopes, Apartado de Correos 321, E-38700 Santa Cruz de Palma, Spain", "School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews, Fife KY16 9SS, UK", "Department of Physical Sciences, The Open University, Milton Keynes MK7 6AA, UK", "Observatoire astronomique de l'Université de Genève, 51 ch. des Maillettes, CH-1290 Sauverny, Switzerland", "Department of Physics, University of Cambridge, J J Thomson Av, Cambridge CB3 0HE, UK", "Observatori Astronòmic de Mallorca, Camí de l'Observatori s/n 07144 Costitx, Mallorca, Spain", "Observatoire astronomique de l'Université de Genève, 51 ch. des Maillettes, CH-1290 Sauverny, Switzerland", "Isaac Newton Group of Telescopes, Apartado de Correos 321, E-38700 Santa Cruz de Palma, Spain", "Astrophysics Group, Keele University, Staffordshire ST5 5BG, UK", "Physics and Astronomy Department, Vanderbilt University, Nashville, TN 37235, USA; Department of Physics, Fisk University, Nashville, TN 37208, USA", "Observatoire astronomique de l'Université de Genève, 51 ch. des Maillettes, CH-1290 Sauverny, Switzerland", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK", "Department of Physics, University of Warwick, Coventry CV4 7AL, UK" ]	[ "2014A&A...571A..37S", "2015A&A...579A.136M", "2015MNRAS.446.1478B", "2015RAA....15..117S", "2016A&A...588A.127C", "2016ApJ...823...29A", "2017AstBu..72...67G", "2017IAUS..328...54H", "2018Natur.553..477B", "2018exha.book.....P", "2020A&A...644A..77M", "2021ApJ...916L...1A", "2022A&A...658A..75H", "2022AJ....164...26H", "2022PASP..134h2001A", "2023A&A...674A.120A", "2023A&A...675A..62J", "2023MNRAS.520.1642S", "2024AJ....167...48M" ]	[ "astronomy" ]	6	[ "techniques: photometric", "techniques: radial velocities", "stars: individual: WASP-13", "stars: individual: WASP-32", "planetary systems", "Astrophysics - Earth and Planetary Astrophysics" ]	[ "1962AJ.....67..591K", "1962P&SS....9..719L", "1976Ap&SS..39..447L", "1977A&A....57..383Z", "1980ApJ...241..425G", "1982ApJ...263..835S", "1986ApJ...302..757H", "1996Natur.380..606L", "1997ApJ...485..319S", "2004A&A...428.1001C", "2005ApJ...622.1118O", "2007ApJ...655..550G", "2007MNRAS.375..951C", "2008SPIE.7014E..0JP", "2009A&A...502..391S", "2009A&A...505..853B", "2009ApJ...703L..99W", "2010A&A...517L...1Q", "2010A&A...524A..25T", "2010ApJ...718L.145W", "2010ApJ...722L.224B", "2010MNRAS.401.1505B", "2010MNRAS.408.1606W", "2010MNRAS.408.1666S", "2010PASP..122.1465M", "2011A&A...527L..11H", "2011A&A...533A.113M", "2011A&A...533A.130H", "2011A&A...534L...6T", "2011ApJ...730L..31H", "2011ApJ...741L...1W", "2011ApJ...743...61S", "2011ApJS..197...10B", "2011ApJS..197...14D", "2011MNRAS.412.2790L", "2011MNRAS.413L..71W", "2011MNRAS.414.3023S", "2012ApJ...757...18A", "2012ApJ...759L..36H", "2012ApJ...760..139B", "2013A&A...551A..80T", "2013ApJ...768...79G", "2013ApJ...771...11A", "2013ApJ...774...53B", "2013ApJ...775...54S", "2013ApJ...776L..35Z", "2013MNRAS.436..898K", "2013MNRAS.436.1883W", "2013Sci...342..331H", "2014ApJ...785..116L", "2014MNRAS.438L..31G" ]	[ "10.1093/mnras/stu520", "10.48550/arXiv.1403.4095" ]	1403	1403.4095_arXiv.txt	The study of gas giants orbiting close to their host stars allows an insight into the formation and evolution of exoplanets. For example, combined planetary transit photometry and radial velocity (RV) measurements enables the planetary density to be found, providing constraints on the planetary composition. Whilst this provides clues to the formation processes at work, the Rossiter-McLaughlin (RM) effect is thought to be a complementary probe of exoplanet dynamical histories. The RM effect is measured using in-transit spectroscopic observations, revealing a deviation from the Keplerian orbital motion as the star orbits the barycentre of the star-planet system. The effect is caused by the planet occulting the rotating stellar surface. This introduces an asymmetry in the stellar absorption profile, resulting in an apparent shift of the spectral lines. The RM waveform allows the sky projected spin-orbit alignment angle ($\lambda$) between the rotation axis of the host star and the normal to the planetary orbital plane to be determined. The alignment angle is thought to provide a window on the dynamical evolution of exoplanets. Hot-Jupiters are thought to form beyond the snow-line where icy cores become massive enough to accrete a gaseous envelope before subsequently migrating either via planet-disk, planet-planet or planet-star interactions. Planet-disk interactions are thought to be dynamically gentle \citep{Goldreich1980, Lin1996} and do not peturb the original inclination of the planet. Other migration mechanisms such as planet-planet and planet-star interactions via the Kozai-Lidov mechanism are more dynamically violent \citep{Kozai1962, Lidov1962}. The presence of a third body in the system excites periodic oscillations in the eccentricity and inclination of the orbit, where tidal dissipation and circularisation shrinks the semi-major axis. The oscillating inclination resulting from Kozai-Lidov interactions produces a continuum of inclinations with stable orbits. Thus, it is expected that hot-Jupiters will exhibit misaligned orbits if such migration processes are operating. However, it should be noted that measuring a spin-orbit alignment angle of $\lambda=0^{\circ}$ does not necessarily indicate an aligned planetary system. When the impact parameter is low, the RM waveform is independent of $\lambda$ and instead controls the amplitude, leading to a strong degeneracy between $v \sin i$ and $\lambda$ \citep{Gaudi2007}. For example, in a system with an impact parameter of $b=0$ and/or where the stellar rotation axis is inclined in the direction of the observer, any orientation of the planetary orbit leads to a symmetric RM waveform. By calculating the inclination of the stellar rotation axis, these degeneracies can be broken and the true `3D' system geometry ascertained. Currently, 76\footnote{Holt-Rossiter-McLaughlin Encyclopaedia: \newline http://www.physics.mcmaster.ca/~rheller/} planets have a measured $\lambda$ where 45$\%$ of planets show substantial misalignments. This population of misaligned planets appears to be synonymous with hotter host stars ($T_\mathrm{eff} \geq 6250$K) whilst aligned planets are preferentially observed orbiting cool host stars. One proposed reason for the alignment-misalignment transition is a change in the internal structure of main-sequence host stars around 6250K, where the outer convective envelope is responsible for tidal interactions. Another correlation in current RM data is the degree of alignment with system age \citep{Triaud2011angle}. All systems with M$_{\star} \geq 1.2$ M$_{\odot}$ were considered and systems with ages greater than 2.5 Gyrs are preferentially aligned, whereas those below this age are preferentially misaligned. This reflects the development of the convective envelope with system age and lends further support to alignment arising from tidal interactions. \citet{Albrecht2012} showed that other correlations of alignment with the orbital period, ratio of planet mass to stellar mass and possibly orbital distance with $\lambda$ provide further evidence that realignment is driven by tidal interactions. In order to interpret the results of RM observations as a tracer of dynamical evolution alone, it must be assumed that the original protoplanetary disk and the star are well-aligned. While this seems valid based on angular momentum conservation, theoretical models have begun to challenge this assumption, showing that star-disk misalignment is possible in the pre-mainsequence phase \citep{Bate2010, Lai2011}. Thus, measuring $\lambda$ may not trace planet migration mechanisms but perhaps traces star formation processes or a combination of both. \citet{Watson2011} studied the inclination of resolved debris disks and the inclination of their host stars for 9 systems, showing that all are consistent with alignment. The authors note that all systems have $T_\mathrm{eff} < 6250$K and other candidates with $T_\mathrm{eff} > 6250$K would be important in exploring the full alignment-misalignment theoretical picture proposed by \citet{Winn2010}. Further systems, with a range of spectral types, were investigated by \citet{Greaves2013} where the stellar inclination was found to be aligned with the spatially resolved debris disk for all systems. Recently \citet{Kennedy2013} tested the alignment of the full star-disk-planet system for HD 82943, the first time the full alignment of a system has been investigated. The complete system (the inclination of the stellar rotation axis, the normal to the disk plane and the normal to the planetary orbit) was found to be aligned at a level similar to the Solar System. Another approach to distinguish between primordial star-disk misalignments and misalignment driven by migration is to consider the growing number of multiplanet systems. \citet{Albrecht2013} recently analysed the multiple-transiting systems KOI-94 \citep{Hirano2012} and Kepler-25, finding $\lambda=-11 \pm 11^{\circ}$ and $\lambda=7 \pm 8^{\circ}$, results consistent with alignment. Whilst this was thought to hint that multi-planet systems migrate via planet-disk interactions and hot-Jupiters migrate by a different pathway, evidence for misaligned multi-planet systems has been found \citep{Huber2013, Walkowicz2013}. It is clear that a full picture of hot-Jupiter formation and migration is far from complete, requiring the continual buidling of statistics, preferably beyond the $T_\mathrm{eff}$ dependence, to explore unstudied regions of parameter space. In this paper we report Rossiter-McLaughlin (RM) observations of WASP-13 and WASP-32. WASP-13 and WASP-32 are both slow rotators \citep{Skillen2009, Maxted2010} and cool stars with effective temperatures $\sim$6000K. Section 2 outlines the observations and analysis procedure. In Section 3 the derived parameters are presented and discussed. Next a search for the stellar rotation period for both systems was investigated. For WASP-32, where a period was found, we then computed the true 3D alignment angle. Finally, we conclude with a discussion of our results in Section 4. \begin{table} \centering \caption{Adopted system parameters and uncertainties used to model the RM effect, and other photometric parameters used in this work. The reference is indicated at the end of the column for each object.} \label{Adopted} \begin{tabular}{lccc} \hline\hline Parameter (units) &Symbol & {WASP-13} & {WASP-32} \\ \hline Orbital Period (days) &$P$ &$~4.3530135\pm~0.0000027$ &$~2.718659\pm~0.000008$ \\[4pt] Transit epoch &$T_\mathrm{0}$ &$2455305.62823 \pm~0.00025~(\mathrm{BJD_{UTC}})$ & $2455151.0546\pm~0.0005~(\mathrm{HJD})$ \\[4pt] Transit duration (hours) &$T_{d}$ &$4.003\,\pm~0.024$ &$2.424\pm~0.048$ \\[4pt] Orbital inclination ($^{\circ}$) & $i$ &$85.43\,\pm~0.29$ &$85.3\,\pm~0.5$ \\[4pt] Planet/Star radius ratio & $R_{p}/R_{}$ &$0.0919\pm~0.0126$ &$0.11\,\pm~0.01$ \\[4pt] Scaled semi-major axis &$a/R_{}$ &$7.54\pm~0.27$ &$7.63\,\pm~0.35$ \\[4pt] Eccentricity &$e$ &$0\,\mathrm{(adopted)}$ &$0.0180\pm~0.0065$ \\[4pt] \hline Reference & & \citet{Chew2013} & \citet{Maxted2010} \\ \hline \end{tabular} \end{table}	\label{Conc} The spectrosopic transits of WASP-13b and WASP-32b were observed with the SOPHIE spectrograph and the projected spin-orbit alignment was determined for both systems where $\lambda={8^{\circ}}^{+13}_{-12}$ and ${-2^{\circ}}^{+17}_{-19}$, respectively. WASP-13 and WASP-32 are consistent with alignment within 1$\sigma$. This suggests WASP-13 and WASP-32 had a gentle migration history and remained unperturbed from the original obliquity of the protoplanetary disk. An alternative scenario is the gradual loss of orbital energy by the planet through tidal dissipation, acting to realign the stellar spin and planetary orbital axes over a long enough timescale \citep{Winn2010}. The misalignment angle measured for WASP-32 in this work is consistent with the value of $\lambda={10.5^{\circ}}^{+6.4}_{-6.5}$ measured by \citet{Brown2012}. Our measured 3D alignment angle of $\psi=11^{\circ} \pm 14$ provides further evidence that the system is well-aligned. It is important to note that $\psi$ has not been measured for many systems (see Table \ref{psi}) and WASP-32 adds to the number of systems with a measured 3D alignment angle. Further, Table \ref{psi} shows that some systems are unambiguously aligned. Attempts have been made to derive the original obliquity distribution (Triaud et al., 2010; Li 2013) assuming a $\mathrm{cos} i_{\star}$ probability distribution in the stellar inclination, however this deprojection technique means that cases where $i_{\star}=90^{\circ}$ are unaccounted for. All current measurements of $\psi$ in Table \ref{psi} show that there is a bimodal distribution in $\psi$: a planetary population that is aligned and one that is near isotropic. Thus, any attempt to deproject the population of spin-orbit angles is destined to fail if current trends in $i_{\star}$ are not recognised. \citet{Winn2010} proposed one mechanism that could explain the observed distribution of alignment angles via tidal dissipation with the host star. In this scheme aligned planets are expected around cool stars ($T_{\mathrm{eff}}<6250$K) and misaligned planets are expected orbiting hot host stars ($T_{\mathrm{eff}}>6250$K). WASP-13 and WASP-32 have $T_{\mathrm{eff}}=5989 \pm 48$K and $T_{\mathrm{eff}}=6100 \pm 100$K, respectively, and therefore both lie in the cool regime. Thus both WASP-13 and WASP-32 add further evidence to alignment arising from tidal interactions. Also it must be noted that WASP-32 has a $T_{\mathrm{eff}}$ close to 6250K, perhaps indicating that it is possible for massive planets (in this case with a mass $>3 \mathrm{M_{J}}$) to tidally realign around relatively hot host stars. Alignment is expected to be determined by planet-star tidal interactions. The tidal interaction timescale due to tidal dissipation in the convective envelope is related to $q$, the planet to star mass ratio, and the scaled semi-major axis, $a/R_{\star}$ (see equation 2 of \citealt{Albrecht2013}): \begin{equation} \label{tidal} \frac{1}{\tau_\mathrm{CE}}=\frac{1}{10 \times 10^{9} \mathrm{yr}}q^2\left(\frac{a/R_{\star}}{40}\right)^{-6} \end{equation} \noindent Thus, the above equation shows that $\tau_\mathrm{CE} \propto q^{-2} \times (a/R_{\star})^6$. As planet-star tidal interactions with the convective envelope are thought to be responsible for aligning hot-Jupiters via tidal dampening, we modified Equation \ref{tidal} to include the convective mass of the planet host, $M_{\mathrm{conv}}$. Thus in Figure \ref{tides} an ensemble of systems with RM measurements are plotted against $(M_p/M_{\mathrm{conv}})^{-1/3} \times (a/R_{\star})$, a quantity proportional to the tidal dissipation timescale. The convective envelope mass, $M_{\mathrm{conv}}$, was derived using the EZ-Web stellar evolution code \footnote{EZ-Web Stellar Evolution Code: \newline http://www.astro.wisc.edu/~townsend/static.php?ref=ez-web}. To run the stellar models, the age of the system is required. For systems lacking derived ages, we have assumed an age of 4 Gyrs, but note the results are largely insensitive to age. Note the $x$-axis scale was chosen such that $\tau_\mathrm{CE}^{1/6} \propto (M_p/M_{\mathrm{conv}})^{-1/3} \times (a/R_{\star})$ for plotting convenience. As suggested by \citet{Zahn1977} tides are dissipated more effectively when the planet orbits a star with a convective envelope. \citet{Winn2010} have postulated that tides have changed the distribution of spin-orbit angles with $T_{\mathrm{eff}}<6150$K but left the distribution unaltered with $T_{\mathrm{eff}}>6350$K. Thus, only `cool' systems with $T_{\mathrm{eff}}<6150$K are plotted in Figure \ref{tides}. It can be observed that as the tides become weaker (when the tidal dissipation timescale increases) there is some evidence that misaligned orbits are more likely. WASP-13 and WASP-32 are plotted in Figure \ref{tides} and are consistent with alignment. A recent addition to the ensemble of RM measurements is WASP-80b \citep{Triaud201380}, a K7-M0 star and the coolest host star in the sample with $T_{\mathrm{eff}}=4145 \pm 100$K. Even as the coolest system, the planet is on an inclined circular orbit with $\|\lambda\|=75^{\circ}$ similar to the spin-orbit angle measured around hotter mid-F stars. This suggests that hot-Jupiters may have been more frequently misaligned in the past. However, other mechanisms could act to misalign a system such as the presence of another peturbing body or if the host star is not old enough to develop a convective envelope. WASP-80b is considered a rare example of a misaligned system around a cool host star \citep{Triaud201380}. However, Figure \ref{tides} suggests that WASP-80b is yet to realign because of its long tidal dissipation timescale. Even though the above analysis is simplified, Figure \ref{tides} suggests that planet-star tidal interactions likely play a role in damping the obliquities of hot-Jupiters around cool host stars. Systems with short tidal dissipation timescales are preferentially aligned, however those with longer timescales show an apparent random distribution in $\lambda$. This may suggest that hot-Jupiters once had a broader range of obliquities in the past and, that they have been realigned over time via tidal interactions \citep{Albrecht2013}. In Figure \ref{tides}, WASP-8b is the most obvious outlier in the distribution, however WASP-8 is a dynamically complex system with suggestions the Kozai mechanism or violent dynamical interations may explain the misaligned orbit \citep{Queloz2010}. \begin{table} \caption{Table of all cases where the 3D alignment angle, $\psi$, has been reported in the literature. The measured $\psi$ and reference is indicated in the table. Multiple references indicate where $\psi$ has been measured in seperate studies. Cases where multiple $\psi$ measurements are listed with a single reference stems from orbital geometry degeneracies. Our result for WASP-32 adds to the number of systems with a complete 3D alignment angle determination.} \label{psi} \begin{tabular}{lcc} \hline\hline Object & $\psi~ (^{\circ})$ & Reference\\ \hline CoRoT-18b & $20 \pm 20$ & [1] \\ HAT-P-7 & $>86.1$ & [2] \\ HAT-P-11 & $106^{+15}_{-11},~ 97^{+8}_{-4}$ & [3] \\ Kepler-16(AB)b &$<18.3$ & [4] \\ Kepler-17b & $0 \pm 15 $ & [5] \\ Kepler-63b & $104^{+9}_{-14}$ & [6] \\ Kepler-13.01 & $54 \pm 4,~56 \pm 4,~124 \pm 4,~126 \pm 4$ & [7] \\ KOI-368.01 & $69^{+9}_{-10}$ & [8] \\ PTFO 8-86956b & $73.1 \pm 0.5$ & [9] \\ WASP-15b & $>90.3$ & [10] \\ WASP-17b & $>91.7,~>92.6$ & [10], [11] \\ WASP-19b & $<19,~<20$ & [10], [12] \\ \textbf{WASP-32b} & $\mathbf{11 \pm 14}$ & \textbf{This Work} \\ \hline \hline \end{tabular} \medskip \textbf{[1]} \citet{Hebrard2011psi} \textbf{[2]} \citet{Winn2009} \textbf{[3]} \citet{Ojeda2011} \textbf{[4]} \citet{Winn2011psi} \textbf{[5]} \citet{Desert2011} \textbf{[6]} \citet{Ojeda2013} \textbf{[7]} \citet{Barnes2011} \textbf{[8]} \citet{Zhou2013} \textbf{[9]} \citet{Barnes2013} \textbf{[10]} \citet{Triaud2010psi} \textbf{[11]} \citet{Bayliss2010} \textbf{[12]} \citet{Hellier2011} \end{table} It is known that stars with $\mathrm{M}>1.2\mathrm{M}_{\odot}$ cool as they evolve along the main sequence. As the star cools an outer convective envelope develops, increasing the strength of the tidal interactions. Thus the distribution of spin-orbit angles is expected to change with time where a planet originally on a misaligned orbit will realign as the convective envelope of the host star develops. Triaud 2011 plotted $\|\lambda\|$ against age for all systems with $\mathrm{M}>1.2\mathrm{M}_{\odot}$. The plot provides weak evidence that the spin-orbit alignment distribution changes with time and is another manifestation of the influence of tidal interactions. Objects with ages 2.5-3 Gyrs appear aligned, whereas more misaligned systems are observed around stars with younger ages. Even though the plots of $\|\lambda\|$ against $a/R_{\star}$ and age show evidence for evolution due to tides, it is still unclear if an original misaligned hot-Jupiter population would survive realignment around `cool' host stars or tidally infall into the star, leaving the aligned population observed today. We have presented RM measurements for WASP-13 and WASP-32. Analysing out-of-transit survey photometry for WASP-32 revealed the rotation period of the host star, and thus the 3D alignment angle $\psi=11^{\circ} \pm 14$ of the planetary system. WASP-32 adds to the number of systems with a full 3D alignment angle determination. It is clear that it is becoming increasingly important to investigate the full star-planet-disk \citep[e.g.,][]{Watson2011, Kennedy2013} alignment in order to fully assess the migration history of exoplanets. Only with an alignment determination of the whole system can we begin to fully evaluate the migration scenarios of hot-Jupiters.	14	3	1403.4095	We present Rossiter-McLaughlin observations of WASP-13b and WASP-32b and determine the sky-projected angle between the normal of the planetary orbit and the stellar rotation axis (λ). WASP-13b and WASP-32b both have prograde orbits and are consistent with alignment with measured sky-projected angles of λ =8°^{+13}_{-12} and λ =-2°^{+17}_{-19}, respectively. Both WASP-13 and WASP-32 have T<SUB>eff</SUB> < 6250 K, and therefore, these systems support the general trend that aligned planetary systems are preferentially found orbiting cool host stars. A Lomb-Scargle periodogram analysis was carried out on archival SuperWASP data for both systems. A statistically significant stellar rotation period detection (above 99.9 per cent confidence) was identified for the WASP-32 system with P<SUB>rot</SUB> = 11.6 ± 1.0 days. This rotation period is in agreement with the predicted stellar rotation period calculated from the stellar radius, R<SUB></SUB>, and vsin i if a stellar inclination of i<SUB></SUB> = 90° is assumed. With the determined rotation period, the true 3D angle between the stellar rotation axis and the planetary orbit, ψ, was found to be ψ = 11° ± 14°. We conclude with a discussion on the alignment of systems around cool host stars with T<SUB>eff</SUB> < 6150 K by calculating the tidal dissipation time-scale. We find that systems with short tidal dissipation time-scales are preferentially aligned and systems with long tidal dissipation time-scales have a broad range of obliquities.	false	[ "systems", "cool host stars", "long tidal dissipation time-scales", "short tidal dissipation time-scales", "prograde orbits", "obliquities", "archival SuperWASP data", "the predicted stellar rotation period", "the stellar rotation axis", "ψ", "the WASP-32 system", "the tidal dissipation time-scale", "measured sky-projected angles", "both systems", "these systems", "A statistically significant stellar rotation period detection", "T<SUB", "the planetary orbit", "rot</SUB", "a stellar inclination" ]	6.643582	14.401163	-1
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P., Jr.", "Campista, M.", "Amorim, A.", "Asari, N. V.", "Bongiovanni, A.", "Bonoli, S.", "Bruzual, G.", "Cardiel, N.", "Cava, A.", "Cid Fernandes, R.", "Coelho, P.", "Cortesi, A.", "Delgado, R. G.", "Diaz Garcia, L.", "Espinosa, J. M. R.", "Galliano, E.", "Gonzalez-Serrano, J. I.", "Falcon-Barroso, J.", "Fritz, J.", "Fernandes, C.", "Gorgas, J.", "Hoyos, C.", "Jimenez-Teja, Y.", "Lopez-Aguerri, J. A.", "Lopez-San Juan, C.", "Mateus, A.", "Molino, A.", "Novais, P.", "OMill, A.", "Oteo, I.", "Perez-Gonzalez, P. G.", "Poggianti, B.", "Proctor, R.", "Ricciardelli, E.", "Sanchez-Blazquez, P.", "Storchi-Bergmann, T.", "Telles, E.", "Schoennell, W.", "Trujillo, N.", "Vazdekis, A.", "Viironen, K.", "Daflon, S.", "Aparicio-Villegas, T.", "Rocha, D.", "Ribeiro, T.", "Borges, M.", "Martins, S. L.", "Marcolino, W.", "Martinez-Delgado, D.", "Perez-Torres, M. A.", "Siffert, B. B.", "Calvao, M. O.", "Sako, M.", "Kessler, R.", "Alvarez-Candal, A.", "De Pra, M.", "Roig, F.", "Lazzaro, D.", "Gorosabel, J.", "Lopes de Oliveira, R.", "Lima-Neto, G. B.", "Irwin, J.", "Liu, J. F.", "Alvarez, E.", "Balmes, I.", "Chueca, S.", "Costa-Duarte, M. V.", "da Costa, A. A.", "Dantas, M. L. L.", "Diaz, A. Y.", "Fabregat, J.", "Ferrari, F.", "Gavela, B.", "Gracia, S. G.", "Gruel, N.", "Gutierrez, J. L. L.", "Guzman, R.", "Hernandez-Fernandez, J. D.", "Herranz, D.", "Hurtado-Gil, L.", "Jablonsky, F.", "Laporte, R.", "Le Tiran, L. L.", "Licandro, J", "Lima, M.", "Martin, E.", "Martinez, V.", "Montero, J. J. C.", "Penteado, P.", "Pereira, C. B.", "Peris, V.", "Quilis, V.", "Sanchez-Portal, M.", "Soja, A. C.", "Solano, E.", "Torra, J.", "Valdivielso, L." ]	2014arXiv1403.5237B	[ "J-PAS: The Javalambre-Physics of the Accelerated Universe Astrophysical Survey" ]	340	[ "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-" ]	[ "2012arXiv1205.4688L", "2014A&A...564A.127L", "2014A&A...571A..60S", 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"2012JCAP...09..009W", "2012MNRAS.419.2133H", "2012MNRAS.420...61K", "2012MNRAS.420..926F", "2012MNRAS.420.1167A", "2012MNRAS.420.1384S", "2012MNRAS.420.1481V", "2012MNRAS.420.1916S", "2012MNRAS.420.2042A", "2012MNRAS.420.2377M", "2012MNRAS.421..872C", "2012MNRAS.421.1569B", "2012MNRAS.421.2355H", "2012MNRAS.421.2904H", "2012MNRAS.422.2187M", "2012MNRAS.422.2904G", "2012MNRAS.423.2308Z", "2012MNRAS.423.3163K", "2012MNRAS.423.3251A", "2012MNRAS.424..157V", "2012MNRAS.424..172R", "2012MNRAS.424..224H", "2012MNRAS.424.1268C", "2012MNRAS.424.1614O", "2012MNRAS.424.2757M", "2012MNRAS.425.1042J", "2012MNRAS.425.2443K", "2012MNRAS.426..549S", "2012MNRAS.426.1767P", "2012MNRAS.426.2046A", "2012MNRAS.426.2142L", "2012MNRAS.426.2566M", "2012MNRAS.427.2079C", "2012MNRAS.427.2132P", "2012MNRAS.427.2146X", "2012MNRAS.427.2168M", "2012MNRAS.427.3044S", "2012MNRAS.427.3435A", "2012MNRAS.427L..35K", "2012PASA...29..359K", "2012PDU.....1...50K", "2012PhLB..716..165A", "2012PhRvD..86h3006S", "2012PhRvD..86h3504Y", "2012PhRvD..86j3513H", "2012PhRvL.109d1101H", "2012SPIE.8446E..0TD", "2012SPIE.8448E..1AC", "2012SPIE.8448E..1VG", "2012SPIE.8451E..16C", "2013A&A...549A..60C", "2013A&A...550A..58V", "2013A&A...553A..78L", "2013A&A...554L...3O", "2013A&A...558A..89K", "2013A&A...558A..90K", "2013AJ....145..138B", "2013ASSL..396..223D", "2013ApJ...762...38P", "2013ApJ...762...43K", "2013ApJ...762L..30Z", "2013ApJ...763...88C", "2013ApJ...763L..41B", "2013ApJ...764...48K", "2013ApJ...765...67M", "2013ApJ...765L...2B", "2013ApJ...766...32B", "2013ApJ...768...29V", "2013ApJ...769...13U", "2013ApJ...769...25S", "2013ApJ...769...91B", "2013ApJ...769..147L", "2013ApJ...770L..15Z", "2013ApJ...771...89O", "2013ApJ...772...63W", "2013ApJ...772...65C", "2013ApJ...779...53Y", "2013ApJ...779..127C", "2013ApJS..204....5K", "2013LRR....16....6A", "2013MNRAS.428..109S", "2013MNRAS.428..778O", "2013MNRAS.428.1827T", "2013MNRAS.428.2885D", "2013MNRAS.428.2980R", "2013MNRAS.429..556M", "2013MNRAS.429.2183P", "2013MNRAS.429.2643C", "2013MNRAS.429.2858M", "2013MNRAS.429L..84K", "2013MNRAS.430..330H", "2013MNRAS.430.2513H", "2013MNRAS.430.2638M", "2013MNRAS.431.3307S", "2013MNRAS.431.3373H", "2013MNRAS.432..318A", "2013MNRAS.432.1133M", "2013MNRAS.432.3141C", "2013MNRAS.433.2706O", "2013MNRAS.433.2857M", "2013MNRAS.434.1443X", "2013MNRAS.434.1604Z", "2013MNRAS.434.2856B", "2013MNRAS.435..158O", "2013MNRAS.435..952S", "2013MNRAS.435.3206D", "2013PhR...530...87W", "2013arXiv1308.0847L", "2013hsa7.conf..115A", "2013hsa7.conf..405S", "2013hsa7.conf..862C", "2014A&A...561A..71Z", "2014A&A...567A.111M", "2014A&A...571A...1P", "2014A&A...571A..17P", "2014A&A...571A..22P", "2014A&A...571A..23P", "2014ApJ...784L..25C", "2014ApJ...794..156R", "2014MNRAS.437.2594W", "2014MNRAS.439...83A", "2014MNRAS.439.1337O", "2014MNRAS.439L..21K", "2014MNRAS.441.1783A", "2014MNRAS.441.2891M", "2014MNRAS.442..589A", "2014MNRAS.442.2680G", "2014PASJ...66R...1T", "2019ApJ...873..111I" ]	[ "10.48550/arXiv.1403.5237" ]	1403	1403.5237_arXiv.txt	The last decade has seen an accumulation of very large field Astrophysical Surveys (area $>5000 \sq\degr$). A key factor in this development has been the undoubted success of the Sloan Digital Sky Survey which has spawned significant advances in almost all the fields in Astrophysics. The quest for the origin of Dark Energy has also been a powerful motivator, fostering many of the current projects like Pan-STARRS \citep{2002SPIE.4836..154K}, DES \citep{des}, and BOSS \citep{BOSS}, and also being one of the main goals of the very large extragalactic surveys planned to start around the beginning of the next decade LSST \citep{lsst}, Euclid \citep{euclid} and DESI \citep{levi2013}. All these surveys are based on two traditional, one century-old, astronomical methods: broad-band imaging ($R \sim 6$)\footnote{For imaging, we define the wavelength resolution as $R_\lambda=\lambda/\Delta_\lambda$, where $\Delta_\lambda$ is the filter width. Another alternative definition would be $R_z=(1+z)/\delta_z$, the inverse of the redshift error. This is usually much higher than $R$ for photometric redshifts, for instance, for broadband imaging $R_z\sim 25$, for J-PAS-like Narrow Band(NB) imaging $R_z\sim 333$} supplemented by moderate resolution spectroscopy ($R \sim 500$). Optical broad-band imaging with traditional astronomical filter systems is observationally efficient but yields very limited redshift information, with, typically, $dz/(1+z) \gtrsim 3\%$, see e.g. \citet{2010A&A...523A..31H} and references therein. Spectroscopy for cosmological purposes provides higher resolution, $dz/(1+z)\sim 0.0005-0.001$ but to be competitive requires very high object multiplexing $\gtrsim 1000$, making state-of-the-art spectrographs extremely expensive and very complex to develop. In addition, the information provided by low-resolution, cosmologically-oriented spectroscopy is relatively limited for other purposes, since the spectra are usually low $S/N$ and for efficiency reasons only objects of direct cosmological interest are systematically targeted. Several projects like COMBO-17 \citep{wolf08}, ALHAMBRA \citep{molino13} and COSMOS \linebreak \citep{ilbert09} have carried out medium band imaging over a few square degrees, hinting at the potential of this approach. These surveys, with $\sim 300\AA$ medium band filters reach precisions of $dz/(1+z)\approx 0.8\%$ and of $0.6\%$ for the highest quality photo-z. This is already not far from the $0.35\%$ precision required for radial BAO measurements. \citet{Benitez2009b} showed that medium band ($R \sim 20$) and narrow band ($R \sim 60$) filter systems are much more effective, in terms of {\it photometric redshift} depth, than what a naive extrapolation from pure {\it photometric} depth would imply. However systematic, multiple-narrow band wide field imaging has not been attempted so far. One objective reason is that until quite recently, it was not possible to build homogeneous filters with a large enough scale. But perhaps the main objection is that NB imaging is quite inefficient for individual objects, since it requires repeated observations to cover a large spectral range; if prompted to consider a NB cosmological survey many would dismiss the idea out of hand \citep{tversky}. However, as explained below, when NB imaging is combined with a large enough FoV, the result is a redshift machine more powerful than any existing spectrograph. Moreover, with a system of contiguous $\sim 100$\AA-width filters it is possible to reach $\approx0.3\%$ redshift precisions for enough LRGs to competitively measure the radial Baryonic Acoustic Oscillation (BAO) scale at $z<1$ \citep{B2009}. J-PAS will observe with an improved version of that system, with the goal of maximizing the effective volume over which we can measure the BAO scale using not only LRGs ($z<1.1$), but also blue galaxies ($z<1.35$) and QSOs ($z<3$), while presenting several features which make the data much more powerful for a wide range of Cosmological and Astrophysical goals. \subsection{{\it Quasi}-Spectroscopy: Wide field Narrow Band Imaging as a Redshift Machine} To understand the power of the J-PAS approach, it is instructive to look at the ``raw'' relative efficiencies of imaging and spectroscopy when observing an object's SED within a particular wavelength range. Let's assume that we have a source with a spectral energy distribution flux $F_\lambda$, observed against a background $B_\lambda$. The imaging system is defined by an average throughput $\eta_I$, a filter width $\Delta\lambda$ and a number of filters $n_f$. The spectrograph is defined by a throughput $\eta_S$. In both cases, we consider a detector with a spatial pixel scale $p_s$ and readout noise $\sigma_{ron}$. The full total exposure time is $t$, divided into $n_r$ individual read outs, and the covered wavelength range is $L=\lambda_{max}-\lambda_{min}$. The central wavelength is thus $\bar\lambda = (\lambda_{max}+\lambda_{min})/2$. For the spectrograph, we have that the signal-to-noise ($S/N$) reached for a fixed time $t$, defined as $q_S=(S/N)_S$ is \begin{equation} q_S=\frac{\bar{F}L\eta_St}{\sqrt{\bar{B}L\eta_St+A_Sn_r\sigma_{ron}^2}} \end{equation} where $\bar{F}$ and $\bar{B}$ are, respectively, the average object flux and background flux in the spectral range $L$ and $A_S$ is the number of pixels covered by the object spectrum on the CCD. For a spectral resolution $R=\lambda/\delta_\lambda$, the pixel scale in the wavelength direction will be $p_\lambda\sim \bar{\lambda}/(2R)$, where we assume at least two pixels per resolution element. If we use a slit of spatial scale $D$, the total number of pixels covered by the spectrum will be $A_S=(L/p_\lambda)\times(D/p_s) \approx 2RD/p_s$ for the typical wavelength range considered here ($L\approx\bar{\lambda}$). For the imaging case, we are using $n_F$ different filters, and we don't cover the whole spectrum in a single shot, only a section $\Delta_\lambda$. Therefore the exposure time for each wavelength segment will be smaller, $t_{exp}=t/n_f$, and $q_I=(S/N)_I$ will be equal to: \begin{equation} q_I=\frac{\bar{F}\Delta_\lambda\eta_It}{\sqrt{\bar{B}\Delta_\lambda\eta_It+A_In_r\sigma_{ron}^2}} \end{equation} Here we assume that $\bar{F}\eta_I \approx \sum(F_i\eta_i)/n_f$ and $\bar{B}\eta_I \approx \sum (B_i\eta_i)/n_f$, where the $i$ values correspond to the individual filters. The total number of pixels covered by the $n_f$ apertures of diameter $D$ is $A_I=n_f\pi(\frac{D}{2p_s})^2$. Let's examine the limit case in which both spectroscopy and imaging are totally background dominated. Then \begin{equation} q_{SB}\approx \frac{\bar{F}}{\sqrt{\bar{B}}}\sqrt{L\eta_St} \end{equation} and \begin{equation} q_{IB}\approx \frac{\bar{F}}{\sqrt{\bar{B}}}\sqrt{\Delta_\lambda\eta_It} \end{equation} Therefore, the relatively $S/N$ ratio of spectroscopy vs. imaging for the background-dominated observation of a single object is \begin{equation} \frac{q_{SB}}{q_{IB}}\approx \sqrt{\frac{L\eta_S}{\Delta_\lambda\eta_I}} \end{equation} If we take as fiducial values $L=9100-3600=5500\AA$, $\Delta_\lambda=145\AA$, $\eta_I=0.7$, $\eta_S=0.25$, we have $\frac{q_{SB}}{q_{IB}}\approx 3.7$\footnote{To make a comparison which focuses on the effects of the filter width, we have not taken into account fiber aperture effects which are not present in imaging, and which in practice would reduce the ratio by a factor of $\sim 2$ in favor of NB imaging}. Thus, as most astronomers will say intuitively, spectroscopy is significantly more efficient than NB imaging (as described here) for a single object observation, since it requires $13.5$ times longer to reach the same $S/N$. The inclusion of readout noise barely changes this result. Let's define, for the imaging case, the ratio $r_I$ between the total readout and background noise \begin{equation} r^2_I=\frac{A_In_r\sigma^2_{ron}}{\bar{B}\Delta_\lambda\eta_It} \end{equation} Then \begin{equation} q_{I}\approx \frac{\bar{F}}{\sqrt{\bar{B}}}\sqrt{\frac{\Delta_\lambda\eta_It}{1+r^2_I}} \end{equation} For the spectroscopic case, since we are assuming the same number of readouts per wavelength segment $n_r$ and the same readout noise $\sigma_{ron}$, we can write \begin{equation} r^2_S=\frac{A_Sn_r\sigma^2_{ron}}{\bar{B}L\eta_St}=\frac{A_s\Delta_\lambda\eta_I}{A_IL\eta_S}r^2_I \end{equation} And therefore \begin{equation} q_S= \frac{\bar{F}}{\sqrt{\bar{B}}}\sqrt{\frac{L\eta_St}{1+\frac{A_S\Delta_\lambda\eta_I}{A_IL\eta_s}r^2_I}} \end{equation} For our fiducial values, the ratio $(L\eta_S)/(\Delta_\lambda\eta_I)\approx 13.5$. The ratio $A_S/A_I=(8Rp_s)/(n_f\pi D)\approx (R_S/100)$, where we have assumed $D=2\arcsec$, $p_s=0.4\arcsec$ and $n_f=54$ Therefore we have \begin{equation} \frac{q_S}{q_I}=\sqrt{\frac{L\eta_S}{\Delta_\lambda\eta_I}}\sqrt{\frac{1+r^2_I}{1+0.07(R_S/100)r^2_I}} \end{equation} In the J-PAS case, assuming $\sigma_{ron}\approx 6e^-$ (as we have done, to be conservative, for all the calculations and mocks presented throughout the paper, although the goal for the camera is $4e^-$), we get $r_I\approx 0.5$. Thus, the inclusion of readout noise for realistic cases does not change things significantly, since the $r_I$-containing factor on the right goes from $1.10$ to $0.85$ for resolutions $R_s=250-4000$. However, the relevant quantity to decide which approach is better as a redshift machine is not the efficiency for an individual object, but the survey speed $v$, which is defined as the total number of objects which can be observed per unit time, with the same signal-to-noise $q$. \begin{equation} v=\frac{N}{t_q} \end{equation} And thus \begin{equation} \frac{v_I}{v_S} \propto \frac{N_It_S}{N_St_I} \end{equation} where $t_I$ and $t_I$ are, respectively, the time required for a spectrograph and the imaging system to reach $S/N=q$. Disregarding the readout noise factors, we have $t_I/t_S=(L\eta_S)/(\Delta_\lambda\eta_I)$, therefore \begin{equation} \frac{v_I}{v_S} \propto \frac{N_I\Delta_\lambda\eta_I}{N_SL\eta_S} \end{equation} The most efficient spectrographs, like BOSS \citep{BOSS}, have $N_S \sim 1000$. In the case of NB imaging, the effective ``multiplexing'' can be extremely large depending on the camera FOV and the density of objects of interest. J-PAS can estimate highly precise photo-z, valid to measure line-of-sight BAOs for $N_I=52,000$ galaxies in $4.7\sq\degr$. Thus, for $<=0.3\%$ photo-z, an instrument like the one which will be used by J-PAS is about 4 times faster, in terms of survey speed, than a 1000-x spectrograph, and it is comparable with a 4000-x spectrograph. \begin{equation} \frac{v_I}{v_S} \approx 4\left(\frac{n_g}{11000 gals /\sq\degr}\right)\left(\frac{FOV}{4.7 \sq\degr}\right)\left(\frac{1000}{N_s}\right) \end{equation} where $n_g$ is the galaxy density per square degree and $FOV$ is the Field of view of the camera in square degrees. We have endeavored to compare both imaging and spectroscopic techniques in a numerically balanced way, however, there are several "hidden variables" favoring imaging techniques which are more difficult to quantify but are nevertheless important to be aware of. We list some of these here: 1. Selection effects in spectroscopy: The necessity of object pre-selection for multi-object spectrographs introduces many unintended biases to a spectroscopic survey. These include effects due to morphology, magnitude and surface density limits all of which have effects on completeness and window function uncertainties which are largely eliminated in an imaging survey. 2. Astrometry: Astrometric errors which can be induced by systematic epoch effects and proper motion uncertainties may lead to small position errors which critically impact aperture coupling efficiencies. Uniform magnitude selection: Spectroscopic optimization of detector real estate requires selection of objects of uniform brightness which can introduce further selection biases. 3. Sky subtraction: Object multiplex has always to be traded with sampling of the sky both for multi-fibre and multi-slit spectroscopy. Sky subtraction for imaging, if done carefully, is generally a much more robust technique which minimally impacts both stochastic and systematic errors. 4. Acceptance aperture: Spectroscopic apertures are always limited by the design of the spectrograph and cannot generally be optimized for particular atmospheric conditions. The size of the instrument scales linearly with aperture; the cost is a much steeper function. There is thus always a strong driver for limiting the size of the aperture which is especially demanding for large telescope spectrographs. For point sources this simply means reduced aperture coupling ratios which are determined by variable seeing conditions, while for marginally resolved sources this leads to incomplete sampling of the intrinsic morphology. 6. Atmospheric dispersion and differential atmospheric refraction: These effects can strongly perturb spectroscopic efficiency as a function of wavelength thus producing errors in the observed source SEDs which are very difficult to impossible to calibrate out. 7. Fibre effects: Spectrograph efficiencies are a function of fibre properties such as throughput, focal ratio degradation and non-telecentric feeds. The first two effects can be a function of fibre placement and telescope orientation (and consequently time) and are very difficult to quantify. Non-telecentric feeds can be well determined but inevitably lead to reductions in efficiency. Many large telescope fibre systems also require multi-fibre connectors which again impact efficiency; it is indeed a fact that very few instrumentation papers actually quote the overall fibre spectrograph efficiency as measured on the sky. Of course there are negatives for multi-band imaging surveys the most prominent of which is the fact that different wavelength samples are taken sequentially which introduce variable efficiencies and PSFs. However, this information can be recovered from the data itself with carefully monitoring of the photometric conditions and by tying the photometry to an all-sky photometric calibration survey, supplied in the case of J-PAS by the T80 telescope. Another limitation is in the effective spectral resolution given by wave-band limitations. However a concerted effort has been made to optimize the band-pass of the J-PAS survey through exhaustive S/N modeling; it is incidental but fortunate that this optimization has led to band-passes which are obtainable through a standard interference filter fabrication processes. But perhaps the main advantage of NB imaging lies in the relative simplicity and low cost (about an order of magnitude lower) of the required instrumentation, specially when compared with a spectrograph with several thousand multiplexing.		14	3	1403.5237	The Javalambre-Physics of the Accelerated Universe Astrophysical Survey (J-PAS) is a narrow band, very wide field Cosmological Survey to be carried out from the Javalambre Observatory in Spain with a purpose-built, dedicated 2.5m telescope and a 4.7 sq.deg. camera with 1.2Gpix. Starting in late 2015, J-PAS will observe 8500sq.deg. of Northern Sky and measure $0.003(1+z)$ photo-z for $9\times10^7$ LRG and ELG galaxies plus several million QSOs, sampling an effective volume of $\sim 14$ Gpc$^3$ up to $z=1.3$ and becoming the first radial BAO experiment to reach Stage IV. J-PAS will detect $7\times 10^5$ galaxy clusters and groups, setting constrains on Dark Energy which rival those obtained from its BAO measurements. Thanks to the superb characteristics of the site (seeing ~0.7 arcsec), J-PAS is expected to obtain a deep, sub-arcsec image of the Northern sky, which combined with its unique photo-z precision will produce one of the most powerful cosmological lensing surveys before the arrival of Euclid. J-PAS unprecedented spectral time domain information will enable a self-contained SN survey that, without the need for external spectroscopic follow-up, will detect, classify and measure $\sigma_z\sim 0.5\%$ redshifts for $\sim 4000$ SNeIa and $\sim 900$ core-collapse SNe. The key to the J-PAS potential is its innovative approach: a contiguous system of 54 filters with $145Å$ width, placed $100Å$ apart over a multi-degree FoV is a powerful "redshift machine", with the survey speed of a 4000 multiplexing low resolution spectrograph, but many times cheaper and much faster to build. The J-PAS camera is equivalent to a 4.7 sq.deg. "IFU" and it will produce a time-resolved, 3D image of the Northern Sky with a very wide range of Astrophysical applications in Galaxy Evolution, the nearby Universe and the study of resolved stellar populations.	false	[ "Stage IV", "resolved stellar populations", "J-PAS", "BAO", "Cosmological Survey", "Dark Energy", "J-PAS unprecedented spectral time domain information", "Northern Sky", "Euclid", "the first radial BAO experiment", "Galaxy Evolution", "Universe", "Astrophysical applications", "first", "~0.7 arcsec", "external spectroscopic follow-up", "FoV", "a multi-degree FoV", "SNeIa", "the Accelerated Universe Astrophysical Survey" ]	12.028275	5.078722	-1
437573	[ "Tapia, Trinidad", "Eliche-Moral, M. Carmen", "Querejeta, Miguel", "Balcells, Marc", "César González-García, A.", "Prieto, Mercedes", "Aguerri, J. Alfonso L.", "Gallego, Jesús", "Zamorano, Jaime", "Rodríguez-Pérez, Cristina", "Borlaff, Alejandro" ]	2014A&A...565A..31T	[ "Evolution induced by dry minor mergers onto fast-rotator S0 galaxies" ]	27	[ "Instituto de Astrofísica de Canarias, C/ Vía Láctea, 38200, La Laguna, Tenerife, Spain ; Departamento de Astrofísica, Universidad de La Laguna, 38200, La Laguna, Tenerife, Spain; Present address: Instituto de Astronomía, Universidad Nacional Autónoma de México, Apdo. 877, Ensenada, BC 22800, Mexico", "Departamento de Astrofísica y CC. de la Atmósfera, Universidad Complutense de Madrid, 28040, Madrid, Spain", "Max-Planck-Institut für Astronomie, Königstuhl, 17, 69117, Heidelberg, Germany", "Isaac Newton Group of Telescopes, Apartado 321, 38700, Santa Cruz de La Palma, Canary Islands, Spain; Instituto de Astrofísica de Canarias, C/ Vía Láctea, 38200, La Laguna, Tenerife, Spain; Departamento de Astrofísica, Universidad de La Laguna, 38200, La Laguna, Tenerife, Spain", "Present address: Instituto de Ciencias del Patrimonio, CSIC, Rúa San Roque 2, Santiago de Compostela, 15704, A Coruña, Spain; Instituto de Astrofísica de Canarias, C/ Vía Láctea, 38200, La Laguna, Tenerife, Spain; Departamento de Astrofísica, Universidad de La Laguna, 38200, La Laguna, Tenerife, Spain", "Instituto de Astrofísica de Canarias, C/ Vía Láctea, 38200, La Laguna, Tenerife, Spain; Departamento de Astrofísica, Universidad de La Laguna, 38200, La Laguna, Tenerife, Spain", "Instituto de Astrofísica de Canarias, C/ Vía Láctea, 38200, La Laguna, Tenerife, Spain; Departamento de Astrofísica, Universidad de La Laguna, 38200, La Laguna, Tenerife, Spain", "Departamento de Astrofísica y CC. de la Atmósfera, Universidad Complutense de Madrid, 28040, Madrid, Spain", "Departamento de Astrofísica y CC. de la Atmósfera, Universidad Complutense de Madrid, 28040, Madrid, Spain", "Departamento de Astrofísica y CC. de la Atmósfera, Universidad Complutense de Madrid, 28040, Madrid, Spain", "Departamento de Astrofísica y CC. de la Atmósfera, Universidad Complutense de Madrid, 28040, Madrid, Spain" ]	[ "2014A&A...570A.103B", "2015A&A...573A..78Q", "2015A&A...579L...2Q", "2015FrASS...2....4D", "2015MNRAS.451.1158P", "2016ASSL..418...15M", "2016MNRAS.463..170C", "2016RMxAA..52...11A", "2017A&A...604A.105T", "2017MNRAS.467.4540B", "2018A&A...609A.132C", "2018A&A...615A..26B", "2018A&A...617A.113E", "2018ApJ...862L..12S", "2018MNRAS.474.1307M", "2018MNRAS.476.3781R", "2018MNRAS.476.4543B", "2019MNRAS.487.4939M", "2019MNRAS.488.4117H", "2020AJ....160...95S", "2020ApJ...888....4S", "2020MNRAS.491.1311M", "2020MNRAS.495.4548B", "2020MNRAS.498.3852B", "2021ApJ...908..135S", "2021MNRAS.506.5030M", "2021MNRAS.507.4262G" ]	[ "astronomy" ]	15	[ "galaxies: bulges", "galaxies: evolution", "galaxies: elliptical and lenticular", "cD", "galaxies: interactions", "galaxies: structure", "galaxies: kinematics and dynamics", "Astrophysics - Galaxy Astrophysics", "Astrophysics - Cosmology and Extragalactic Astrophysics" ]	[ "1978MNRAS.183..501B", "1982ApJ...257...75K", "1982ApJ...257..423H", "1983ApJ...265..632K", "1983ApJ...266...41D", "1983ApJ...266..516D", "1988A&A...193L...7B", "1989ApJS...70..419H", "1990A&A...230...37G", "1991ApJ...383..112F", "1993IAUS..153..209K", "1994ApJ...423..207S", "1995ApJ...447L..87M", "1995MNRAS.277.1341K", "1996AJ....111.2243T", "1996ApJ...471..115B", "1998MNRAS.298..267V", "2000AJ....120..703H", "2001ApJ...563..118P", "2002ApJ...570..610A", "2002ApJ...577..651B", "2004A&A...418L..27B", "2004A&A...423..481K", "2004ApJ...613L..29M", "2004MNRAS.350...35F", "2005A&A...437...69B", "2005ApJ...620L..79S", "2005MNRAS.357..753G", "2005MNRAS.362.1319L", "2005MNRAS.363..937B", "2005MNRAS.364.1105S", "2005RMxAC..23..101K", "2006A&A...457...91E", "2006A&A...458..101A", "2006AJ....131.1336S", "2006ApJ...636L..81N", "2006ApJ...637..214M", "2006ApJ...649L..79M", "2006ApJ...650..791C", "2006ApJS..167....1B", "2006MNRAS.372L..78G", "2006MNRAS.373.1013C", "2007A&A...470..173B", "2007A&A...476.1179B", "2007ApJ...655..790C", "2007ApJ...658..710N", "2007ApJ...660.1151D", "2007ApJ...671.1503M", "2007MNRAS.376..997J", "2007MNRAS.379..401E", "2007MNRAS.379..418C", "2008A&A...477..437D", "2008AN....329.1025B", "2008ApJ...672L.103K", "2008ApJ...685..897B", "2009A&A...496...51P", "2009A&A...496..381H", "2009A&A...501..437Y", "2009A&A...504..389C", "2009A&A...507.1313H", "2009AJ....138..579K", "2009ApJ...692..298W", "2009ApJ...692L..34L", "2009ApJ...697.1290B", "2009ApJ...697L.137P", "2009ApJ...699L.178N", "2009MNRAS.397..802H", "2009MNRAS.397.1202J", "2010A&A...515A...3H", "2010A&A...519A..55E", "2010A&A...521A..71M", "2010AIPC.1240..423T", "2010AJ....140..962M", "2010AJ....140.1814Y", "2010ApJ...708..841W", "2010ApJ...721..259B", "2010ApJ...725..542H", "2010IAUS..262..432T", "2010MNRAS.401.1099H", "2010MNRAS.403.1009M", "2010MNRAS.405.1089L", "2010MNRAS.407.1231R", "2010MNRAS.408.1417S", "2010arXiv1003.0686E", "2011A&A...533A.104E", "2011ApJ...742..103L", "2011MNRAS.410.1197C", "2011MNRAS.410.2625P", "2011MNRAS.412..684B", "2011MNRAS.412L...6B", "2011MNRAS.414..888E", "2011MNRAS.415...32S", "2011MNRAS.415.1783B", "2011MNRAS.416.1654B", "2011MNRAS.416.1680C", "2011MNRAS.417..845K", "2011MNRAS.417..863D", "2012A&A...544A..99W", "2012A&A...547A..48E", "2012A&AT...27..263M", "2012AdAst2012E..28A", "2012ApJ...744...63O", "2012ApJ...744L..11E", "2012ApJ...753...43K", "2012ApJS..198....2K", "2012MNRAS.424.1495L", "2012MNRAS.427..790S", "2013A&A...552A..67E", "2013A&A...553A.116V", "2013ApJ...764..123S", "2013ApJ...771..120P", "2013ApJ...773...34G", "2013MNRAS.428..999P", "2013MNRAS.428.1088M", "2013MNRAS.430.3489L", "2013MNRAS.432..430B", "2013MNRAS.432.1709C", "2014MNRAS.439.1160R", "2014MNRAS.444.3357N" ]	[ "10.1051/0004-6361/201321386", "10.48550/arXiv.1403.2430" ]	1403	1403.2430.txt	\label{sec:introduction} Recent studies have shown that classifying early-type galaxies (ETGs) into fast and slow rotators provides a more consistent distinction in terms of their physical properties than the traditional morphological classification into ellipticals and S0s \citep[E11 hereafter]{Emsellem2007,Emsellem2011}. This is because this criterion is almost independent of the viewing angle (E11), whereas S0 galaxies can be morphologically confused with ellipticals in face-on views \citep[][B11 hereafter]{Bois2011}. According to this classification, the vast majority of ETGs are fast rotators when considered as single-component systems, meaning that they have a noticeable regular rotation pattern, with aligned photometric and kinematic axes, they host inner discs and often bars, and span a wide range of apparent ellipticities ($0<\epse<0.85$). Only a small fraction of ETGs are slow rotators ($\sim$ 15\%), and usually have complex stellar velocity fields and kinematically decoupled cores (E11). Approximately 10-20\% of lenticular galaxies (S0s) in the ATLAS$^\mathrm{3D}$ sample exhibit hybrid properties between fast and slow rotators, lying in the limiting region defined to isolate these two families of objects, with ellipticities spanning the whole range up to $\epse \sim 0.7$ (E11). Fast rotators may be the result of the rebuilding of a stellar disc around a central spheroid, which in turn may come from the destruction of a pre-existing disc through mergers, or by gas exhaustion in spirals \citep[see][]{Khochfar2011}. Observations indicate that these evolutionary mechanisms might depend on environment. Gas stripping and strangulation seem to have been responsible of transforming spirals into S0s in clusters since $z\sim 0.8$ \citep{Barr2007,Desai2007}, while minor mergers, galaxy harassment, and tidal interactions may have triggered an even more dramatic evolution in groups during the same period of time \citep[][]{Moran2007,Bekki2011}. This variety of formation processes agrees pretty well with the diversity of properties exhibited by S0s depending on the environment and the stellar mass \citep{Laurikainen2010,Roche2010,Wei2010,Silchenko2012,Barway2013}. A description of the processes that may have been relevant for the formation of S0s can be found in \citet{Aguerri2012}. Numerical studies indicate that simple fading by itself is not sufficient to produce a fast rotator \citep{Khochfar2011}. However, simulations of major mergers have succeeded in producing both slow- and fast-rotator remnants \citep[][B11]{Gonzalez-Garcia2006a,Jesseit2009}. The fast rotators formed in this way have intermediate apparent flattening ($0.4< \epse <0.6$) and high rotational support ($\lambdae > 0.4$), whereas the resulting slow rotators span the whole range of ellipticities and usually host kinematically decoupled components (B11). Mergers of disc galaxies with higher mass ratios (3:1 and 6:1) basically give rise to fast rotators with intermediate-to-high ellipticities and high rotational support, too (B11). This means that the remnants resulting from mergers with mass ratios lower than 6:1 cannot properly reproduce the region in the $\lambdae - \epse$ parameter space populated by slow rotators with low apparent ellipticities and by galaxies with intermediate properties between fast and slow rotators \citep[with $0.1\lesssim \lambdae \lesssim0.3$, see][B11]{Burkert2008,Khochfar2011}. Intermediate and minor mergers are expected to be much more frequent than major ones in standard hierarchical scenarios \citep{Naab2009,Bezanson2009,Hopkins2010} and more likely to produce S0-like remnants \citep[][]{Bournaud2004,Bournaud2005}. Therefore, it is straightforward to question whether mergers of mass ratios higher than 6:1 are a feasible evolutionary channel for giving rise to S0s with intermediate kinematic properties or not. \citet{Naab2013} have recently shown that subsequent minor mergers in a cosmological context can explain the rare class of slow rotators with low ellipticities. However, their simulations hardly reproduce the location in the \lambdae\ -- \epse\ diagram of the S0s with hybrid kinematic properties ($0.15<\lambdae <0.25$) and $\epse > 0.3$ (see their Fig.\,11). This does not necessarily mean that mergers must be discarded as a feasible mechanism to explain the properties of these galaxies. Cosmological N-body simulations have the advantage (over idealized binary merger simulations) of analysing more realistic pathways to form galaxies, but they are also more limited in numerical resolution. This problem directly affects the way baryons accumulate in the centre of the potential wells, the physics of star formation, the rotational support of the gas component, and the formation of substructures \citep[see][]{Bournaud2008,Piontek2011,Regan2013}. To complement the numerical studies cited above, we have investigated whether dry mergers with mass ratios ranging from 6:1 to 18:1 can explain the formation of S0s with intermediate kinematic properties. We used N-body simulations of binary mergers, starting from gas-poor progenitors with high initial intrinsic ellipticities and rotational support. In this paper, we describe the results obtained for S0s that have spherical original bulges. The effects of considering non-axisymmetric primary bulges will be explored in a forthcoming paper, as the vertical buckling of an original bar (induced by the encounter or by simple natural secular evolution) implies changes in both the velocity ellipsoid and the structure of the bulge \citep{Mihos1995,Martinez-Valpuesta2004,Martinez-Valpuesta2006,Saha2013}. The paper is organized as follows: the models are described in Sect.~\ref{sec:simulations}. Section\,\ref{sec:fastslow} describes the analysis of the global structure and rotational support of the remnants that were produced in our simulations, and compares the results with the distributions of fast and slow rotators obtained in observational surveys and in previous studies of major merger simulations. In Sect.\,\ref{sec:anisotropy} we also analyse the relation between the anisotropy of velocities and the intrinsic ellipticity in our remnants and compare this again with data and previous simulations. Section\,\ref{sec:bulges} shows how different the intrinsic shape and the rotational support of the central remnant bulge can be from those computed for the galaxy as a whole through $\lambdae$ and $\epse$. Section~\ref{sec:relation} shows the relation between the bulge triaxiality and the global rotational support of the whole remnant, two properties that are usually considered to be strongly related. The limitations of the models are commented on Sect.~\ref{sec:limitations}. Finally, the discussion and the main conclusions are provided in Sects.~\ref{sec:discussion} and \ref{sec:conclusions}. \begin{table} \begin{minipage}[t]{\textwidth} \caption{Parameters of the minor and intermediate merger experiments.} \label{tab:models} \centering \begin{tabular}{lccrcc} \hline\hline \multicolumn{1}{c}{Model code} & $M_{\mathrm{sat}}/M_{\mathrm{prim}}$ & $R_{\mathrm{per}}/h_\mathrm{D,prim}$ & \multicolumn{1}{c}{$\theta$} & \multicolumn{1}{c}{\textrm{$(B/D)_\mathrm{prim}$}} & $\alphaTF$ \\ \multicolumn{1}{c}{(1)} & \multicolumn{1}{c}{(2)} & \multicolumn{1}{c}{(3)} & \multicolumn{1}{c}{(4)} & \multicolumn{1}{c}{(5)} & \multicolumn{1}{c}{(6)} \vspace{0.05cm}\\\hline\vspace{-0.3cm}\\ (a)\,\,\, M6 Ps Db & 1:6 (M6) & 0.73 (Ps) & 30 (D) & \textrm{0.5} (b) & 3.5 \\ (a2) M6 Ps Db TF3 & 1:6 (M6) & 0.73 (Ps) & 30 (D) & \textrm{0.5} (b) & 3.0 \\ (a3) M6 Ps Db TF4 & 1:6 (M6) & 0.73 (Ps) & 30 (D) & \textrm{0.5} (b) & 4.0 \\ (b)\,\,\, M6 Ps Rb & 1:6 (M6) & 0.73 (Ps) & 150 (R) & \textrm{0.5} (b) & 3.5 \\ (c)\,\,\, M6 Pl Db & 1:6 (M6) & 8.25 (Pl) & 30 (D) & \textrm{0.5} (b) & 3.5 \\ (d)\,\,\, M6 Pl Rb & 1:6 (M6) & 8.25 (Pl) & 150 (R) & \textrm{0.5} (b) & 3.5 \\ (e)\,\,\, M6 Ps Ds & 1:6 (M6) & 0.87 (Ps) & 30 (D) & \textrm{0.08} (s) & 3.5 \\ (f)\,\,\, M6 Ps Rs & 1:6 (M6) & 0.87 (Ps) & 150 (R) & \textrm{0.08} (s) & 3.5\\ \vspace{-0.4cm}\\\hline\vspace{-0.3cm}\\ (g)\,\,\, M9 Ps Db & 1:9 (M9) & 0.79 (Ps) & 30 (D) & \textrm{0.5} (b) & 3.5 \\ (g2) M9 Ps Db TF3 & 1:9 (M9) & 0.79 (Ps) & 30 (D) & \textrm{0.5} (b) & 3.0 \\ (g3) M9 Ps Db TF4 & 1:9 (M9) & 0.79 (Ps) & 30 (D) & \textrm{0.5} (b) & 4.0 \\ (h)\,\,\, M9 Ps Rb & 1:9 (M9) & 0.79 (Ps) & 150 (R) & \textrm{0.5} (b) & 3.5\\ \vspace{-0.4cm}\\\hline\vspace{-0.3cm}\\ (i)\,\,\, M18 Ps Db & 1:18 (M18)& 0.86 (Ps) & 30 (D) & \textrm{0.5} (b) & 3.5 \\ (j)\,\,\, M18 Ps Rb & 1:18 (M18)& 0.86 (Ps) & 150 (R) & \textrm{0.5} (b) & 3.5 \\ (k)\,\,\,M18 Pl Db & 1:18 (M18)& 8.19 (Pl) & 30 (D) & \textrm{0.5} (b) & 3.5\\ (l)\,\,\, M18 Pl Rb & 1:18 (M18)& 8.19 (Pl) & 150 (R) & \textrm{0.5} (b) & 3.5 \\\hline\\ \end{tabular} \begin{minipage}[t]{\textwidth}{\small \emph{Columns}: (1) Model code: M$m$P[l/s][D/R][b/s][TF3/4], see the text. (2) Luminous mass ratio between the satellite and the primary S0 galaxy. (3) Orbital first pericentre distance in units of the primary disc scale-length. (4) Initial angle between the angular momenta of the orbit and the primary disc. (5) Bulge-to-disc ratio of the original primary S0: $B/D=0.5$ (S0b) or $B/D=0.08$ (S0c). (6) Value of $\alphaTF$ assumed for the scaling of the satellite to the primary S0. More details of the models in EM06 and EM11.} \end{minipage} \end{minipage} \end{table}	\label{sec:conclusions} We have investigated whether minor mergers can explain the existence of S0s with kinematic properties intermediate between fast and slow rotators which major merger and cosmological simulations find difficult to reproduce. We analysed the properties of the remnants that result from dry mergers with mass ratios ranging 6:1 -- 18:1 onto original S0s that initially are fast rotators with high intrinsic ellipticities (they might in turn derive from gas stripping). We found that the minor mergers decrease the intrinsic ellipticity of the whole galaxy, but the remnants do not exhibit a higher percentage of random projections with low \epse\ values. Instead, they increase the fraction of projections with intermediate apparent ellipticities ($0.4<\epse<0.7$) due to the formation of non-axisymmetric distortions and to disc thickening. This means that the remnants become more triaxial. Minor mergers also induce a lower decrease of the rotational support in the remnants than major mergers. These combined effects produce S0 remnants that extend over the limiting region between the distributions of fast and slow rotators in the \lambdae\ -- \epse\ diagram, spanning the whole range of apparent ellipticities up to $\epse \sim 0.8$. Therefore, minor mergers are a plausible mechanism to generate S0s with hybrid kinematics and shape properties. Considering the intrinsic properties of the remnant bulges, we find that the simulated mergers tend to decrease the velocity anisotropy of this sub-component (increasing the rotational support of the bulge or decreasing its intrinsic ellipticity). The remnant bulges remain nearly spherical, but exhibit a wide range of triaxialities ($0.20<T<1.0$). In addition, we showed that the triaxiality of the bulge is only indirectly related with the global rotational support of the whole remnant. In the plane of global anisotropy of velocities ($\delta$) vs.\,intrinsic ellipticity ($\epsilon_\mathrm{e,intr}$), some of our models extend the linear trend found in previous major merger simulations towards higher $\epsilon_\mathrm{e,intr}$ values, while others depart from it, depending on the progenitor. This contributes to increase the dispersion in the diagram. We compared these trends with those exhibited by elliptical and S0 galaxies from the SAURON and ATLAS$^\mathrm{3D}$ projects. While most real ellipticals closely follow the linear trend (consistent with a major merger origin, as already known), S0s are widepreadly in the $\delta$ -- $\epsilon_\mathrm{e,intr}$ diagram. In fact, less than $\sim 40$--50\% of real S0s are located within 1$\sigma$ of the linear trend drawn by major merger simulations. Our simulations show that minor mergers can explain the dispersion exhibited by real S0s in this diagram. Therefore, the different trends exhibited by ellipticals and S0 galaxies in the $\delta$ -- $\epsilon_\mathrm{e}$ diagram probably point to the different role played by major mergers in the build-up of each galaxy type. \begin{figure}[th] \centering \includegraphics[width=4.5cm]{FIGURES/lambda_Emsellem_thickening_M2D.eps} \includegraphics[width=4.5cm]{FIGURES/lambda_Emsellem_thickening_M2R.eps} \includegraphics[width=4.5cm]{FIGURES/lambda_Emsellem_thickening_M3D.eps} \includegraphics[width=4.5cm]{FIGURES/lambda_Emsellem_thickening_M3R.eps} \caption{Dependence on the number of particles used in the simulation of the location in the $\lambda_\mathrm{e}$ -- $\epse$ diagram of the 200 random projections of our models for an identical set of initial conditions (indicated in each frame). \emph{Red dots}: models with $N = 185$K particles. \emph{Green dots}: models with $N = 555$K particles. \emph{Blue dots}: models with $N = 1\,850$K particles. The legend for the lines is the same as in Fig.\,\ref{fig:lambdaobs}.} \label{fig:thickening} \end{figure} % Do not delete the next line \small % Do not delete % %%% Comment the following line if you do not have acknowledgments.	14	3	1403.2430	Context. Numerical studies have shown that the properties of the S0 galaxies with kinematics intermediate between fast and slow rotators are difficult to explain by a scenario of major mergers. <BR /> Aims: We investigate whether the smoother perturbation induced by minor mergers can reproduce these systems. <BR /> Methods: We analysed collisionless N-body simulations of intermediate and minor dry mergers onto S0s to determine the structural and kinematic evolution induced by the encounters. The original primary galaxies represent gas-poor fast-rotator S0b and S0c galaxies with high intrinsic ellipticities. The original bulges are intrinsically spherical and have low rotation. Different mass ratios, parent bulges, density ratios, and orbits were studied. <BR /> Results: Minor mergers induce a lower decrease of the global rotational support (as provided by λ<SUB>e</SUB>) than encounters of lower mass ratios, which results in S0s with properties intermediate between fast and slow rotators. The resulting remnants are intrinsically more triaxial, less flattened, and span the whole range of apparent ellipticities up to ɛ<SUB>e</SUB> ~ 0.8. They do not show lower apparent ellipticities in random projections than initially; on the contrary, the formation of oval distortions and the disc thickening increase the percentage of projections at 0.4 < ɛ<SUB>e</SUB> < 0.7. In the experiments with S0b progenitor galaxies, minor mergers tend to spin up the bulge and to slightly decrease its intrinsic ellipticity, whereas in the cases of primary S0c galaxies they keep the rotational support of the bulge nearly constant and significantly decrease its intrinsic ellipticity. The remnant bulges remain nearly spherical (B/A ~ C/A> 0.9), but exhibit a wide range of triaxialities (0.20 < T < 1.00). In the plane of global anisotropy of velocities (δ) vs. intrinsic ellipticity (ɛ<SUB>e,intr</SUB>), some of our models extend the linear trend found in previous major merger simulations towards higher ɛ<SUB>e,intr</SUB> values, while others clearly depart from it (depending on the progenitor S0). This is consistent with the wide dispersion exhibited by real S0s in this diagram compared with ellipticals, which follow the linear trend drawn by major merger simulations. <BR /> Conclusions: The smoother changes induced by minor mergers can explain the existence of S0s with intermediate kinematic properties between fast and slow rotators that are difficult to explain with major mergers. The different trends exhibited by ellipticals and S0 galaxies in the δ - ɛ<SUB>e</SUB> diagram may be pointing to the different role played by major mergers in the build-up of each morphological type.	false	[ "major merger simulations", "major mergers", "previous major merger simulations", "Minor mergers", "minor mergers", "lt", "high intrinsic ellipticities", "SUB", "intrinsic ellipticity", "lower apparent ellipticities", "S0 galaxies", "S0b progenitor galaxies", "lower mass ratios", "apparent ellipticities", "intermediate kinematic properties", "primary S0c galaxies", "Different mass ratios", "S0c galaxies", "e</SUB", "intermediate and minor dry mergers" ]	11.243109	6.504218	189
812388	[ "Sousa, L.", "Avelino, P. P." ]	2014PhRvD..89h3503S	[ "Stochastic gravitational wave background generated by cosmic string networks: The small-loop regime" ]	21	[ "Rua das Estrelas, 4150-762 Porto, Portugal, Centro de Astrofísica da Universidade do Porto, Rua das Estrelas, 4150-762 Porto, Portugal; , Rua do Campo Alegre 687, 4169-007 Porto, Portugal, Departamento de Física e Astronomia da Faculdade de Ciências da Universidade do Porto, Rua do Campo Alegre 687, 4169-007 Porto, Portugal", "Rua das Estrelas, 4150-762 Porto, Portugal, Centro de Astrofísica da Universidade do Porto, Rua das Estrelas, 4150-762 Porto, Portugal; , Rua do Campo Alegre 687, 4169-007 Porto, Portugal, Departamento de Física e Astronomia da Faculdade de Ciências da Universidade do Porto, Rua do Campo Alegre 687, 4169-007 Porto, Portugal; Sydney Institute for Astronomy, School of Physics, A28, The University of Sydney, New South Wales 2006, Australia" ]	[ "2014IAUS..306..391S", "2015Univ....1....6A", "2016PhRvD..93b3519A", "2016PhRvD..94f3529S", "2018PhRvD..98l3505G", "2020JCAP...04..034A", "2020PhRvD.101j3508S", "2020PhRvD.102j3516F", "2021PhRvD.104b3507R", "2022JCAP...11..024R", "2023EPJC...83.1010W", "2023JHEP...08..196G", "2023LRR....26....5A", "2023PhRvD.108l3516S", "2023arXiv231101300L", "2024JCAP...06..030H", "2024PhRvD.109f3520W", "2024PhRvD.109j3538S", "2024arXiv240309816G", "2024arXiv240413213M", "2024arXiv240503740B" ]	[ "astronomy", "physics" ]	7	[ "98.80.Cq", "Particle-theory and field-theory models of the early Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1976JPhA....9.1387K", "1981PhLB..107...47V", "1984ApJ...282L..51V", "1984Natur.311..109H", "1985NuPhB.252..227K", "1986NuPhB.277..605B", "1988PhRvL..60..257B", "1989NuPhB.319..747A", "1989PhRvD..40..973A", "1990PhRvD..42.2505Q", "1990PhRvL..64..119A", "1992PhRvD..45.1000C", "1992PhRvD..45.3447C", "1995PhRvL..75.4354B", "1996PhRvD..54.2535M", "1996PhRvD..54.7146C", "1997PhRvD..55..573V", "1997PhRvD..56..637V", "1998ApJ...507L.101A", "1998PhRvL..80.2277V", "1998PhRvL..81.2008A", "1998tdc..conf...11B", "2000csot.book.....V", "2001NuPhB.611..125S", "2001PhRvD..64f4008D", "2002PhLB..536..185S", "2002PhRvD..65d3514M", "2002PhRvD..66d3501S", "2004MNRAS.348..105A", "2005APh....23..313C", "2005PhRvD..71f3510D", "2006PhRvD..73d3515M", "2006PhRvD..73j5001S", "2006PhRvD..74d3526H", "2006PhRvD..74f3527V", "2006PhRvD..74h3504P", "2006PhRvL..97b1301S", "2006astro.ph..4069T", "2007JCAP...02..023R", "2007PhRvD..75f3521O", "2007PhRvD..75l3503P", "2007PhRvD..75l5006D", "2007PhRvL..98k1101S", "2008AIPC..983..584M", "2008CQGra..25k4041S", "2008PhRvD..77f3532V", "2008PhRvL.100b1301B", "2009AIPC.1141...10B", "2009PhRvD..80l3523C", "2010CQGra..27h4004K", "2010CQGra..27h4014F", "2010JCAP...10..003L", "2010PhRvD..81j4028O", "2010PhRvD..82f5004B", "2011CQGra..28i4011K", "2011CQGra..28k4002A", "2011PhRvD..83j3507S", "2011PhRvD..84f3502S", "2012JCAP...06..027B", "2012PhRvD..85h3525A", "2012PhRvD..85l2003S", "2012PhRvD..86b3503K", "2013ApJ...762...94D", "2013PhRvD..88b3516S", "2014A&A...571A..16P" ]	[ "10.1103/PhysRevD.89.083503", "10.48550/arXiv.1403.2621" ]	1403	1403.2621_arXiv.txt	} Cosmic strings networks may be produced as a consequence of symmetry-breaking phase transitions \cite{Kibble:1976sj}, being a crucial prediction of many grand-unified scenarios. These networks may survive throughout the cosmological history, potentially leaving behind a variety of observational signatures (see e.g. \cite{Vilenkin:1984ea,Avelino:1997hy,Avelino:1998vu,Sarangi:2002yt,Avelino:2003nn,Bevis:2010gj} and references therein). One such signature is the stochastic gravitational wave background (SGWB) generated by string loops created as a result of string interactions. These loops radiate their energy in gravitational waves (GWs) and their emissions generate a characteristic SGWB \cite{Vilenkin:1981bx,Hogan:1984is,Brandenberger:1986xn,Accetta:1988bg}. The SGWB power spectrum generated by cosmic string networks may be probed using diverse astrophysical experiments: GW detectors \cite{Sigg:2008zz,Accadia:2011zzc,Kuroda:2010zzb,AmaroSeoane:2012km,Kawamura:2011zz}), pulsar timing experiments \cite{Manchester:2007mx,Ferdman:2010xq,Demorest:2012bv}, small-scale fluctuations and B-mode polarization of CMB \cite{Smith:2006nka,Planck:2006aa,Baumann:2008aq}); and big-bang nucleosynthesis \cite{Cyburt:2004yc}. There is thus the prospect either for the detection of specific cosmic string signatures in the SGWB or for the tightening of current constraints on string tension. It is, therefore, important to accurately characterize the SGWB spectrum and to understand its dependence on the large-scale properties of string networks and on the size and emission spectrum of the loops. There are, in the literature, several computations of the SGWB spectrum \cite{Caldwell:1991jj,Caldwell:1996en,Battye:1997ji,Damour:2001bk,Damour:2004kw,Siemens:2006vk,Hogan:2006we,Polchinski:2006ee,DePies:2007bm,Olmez:2010bi,Binetruy:2012ze,Kuroyanagi:2012wm,Sanidas:2012ee,Sousa:2013aaa} based on different assumptions about string network dynamics. In this paper, we present an alternative method to compute the SGWB generated by a realistic cosmic string network and we use it to derive an analytical approximation to the SGWB power spectrum, over a wide frequency range, in the small-loop regime.	} We have proposed an alternative method to compute the SGWB power spectrum generated by cosmic strings, in the small-loop regime. This method does not require underlying simplifications regarding cosmic string network evolution --- avoiding the common assumption of scale-invariance --- thus allowing for an efficient computation of the spectrum generated by string networks undergoing a realistic cosmological evolution. Our method is very useful in the small-loop regime where it is much more efficient than standard methods. This is an important advantage since multiple computations of the spectrum covering a multi-parameter space are often necessary to confront different cosmic string scenarios with the observational data. Moreover, we used this method to derive an analytical approximation to the SGWB spectrum, which was shown to provide an excellent fit to more elaborate results obtained using the VOS model, in the small-loop regime, over a wide frequency range. This analytical approximation constitutes a useful tool for a first estimation of the SGWB power spectrum generated by cosmic string networks, thus allowing for simple estimates of the associated observational constraints in the small-loop regime.	14	3	1403.2621	We consider an alternative approach for the computation of the stochastic gravitational wave background generated by small loops produced throughout the cosmological evolution of cosmic string networks and use it to derive an analytical approximation to the corresponding power spectrum. We show that this approximation produces an excellent fit to more elaborate results obtained using the velocity-dependent one-scale model to describe cosmic string network dynamics, over a wide frequency range, in the small-loop regime.	false	[ "cosmic string network dynamics", "cosmic string networks", "small loops", "the corresponding power spectrum", "a wide frequency range", "the small-loop regime", "the stochastic gravitational wave background", "the cosmological evolution", "an analytical approximation", "more elaborate results", "an excellent fit", "the velocity-dependent one-scale model", "this approximation", "an alternative approach", "the computation", "We", "it", "one" ]	10.382053	-1.727653	86
678201	[ "Chiang, Chi-Ting", "Wagner, Christian", "Schmidt, Fabian", "Komatsu, Eiichiro" ]	2014JCAP...05..048C	[ "Position-dependent power spectrum of the large-scale structure: a novel method to measure the squeezed-limit bispectrum" ]	101	[ "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, Garching, 85741 Germany", "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, Garching, 85741 Germany", "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, Garching, 85741 Germany", "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, Garching, 85741 Germany; Kavli Institute for the Physics and Mathematics of the Universe (Kavli IPMU, WPI), Todai Institutes for Advanced Study, the University of Tokyo, Kashiwa, 277-8583 Japan;" ]	[ "2014JCAP...10..011K", "2014PhRvD..90j3530L", "2014arXiv1412.4671A", "2015JCAP...02..026B", "2015JCAP...08..042W", "2015JCAP...09..028C", "2015JCAP...11..002C", "2015JCAP...11..032G", "2015MNRAS.448L..11W", "2015MNRAS.450.1836K", "2015PhDT........81N", "2015PhRvD..91d3530S", "2015PhRvD..91h3518B", "2015PhRvD..92l3510N", "2015arXiv150803256C", "2016JCAP...06..014T", "2016JCAP...06..025B", "2016PhRvD..93h3517L", "2016PhRvD..94b3002H", "2016PhRvD..94d3519B", "2016PhRvD..94h3528A", "2016PhRvD..94j3506D", "2016PhRvD..94l3502C", "2016arXiv161204247P", "2017JCAP...02..010M", "2017JCAP...06..022C", "2017JCAP...06..042M", "2017JCAP...12..020R", "2017MNRAS.466..780M", "2017MNRAS.466.2496E", "2017MNRAS.470.2100K", "2017MNRAS.471.1581B", "2017PhRvD..95d3529H", "2017PhRvD..95l3512S", "2017PhRvD..95l3517C", "2018JCAP...02..022L", "2018JCAP...07..049C", "2018MNRAS.474.3173P", "2018MNRAS.478.2495N", "2018MNRAS.478.3627A", "2018MNRAS.479..162S", "2018PhR...733....1D", "2018PhRvD..97d3532C", "2018PhRvD..97f3527A", "2018PhRvD..97l3526C", "2018PhRvD..97l3539S", "2018PhRvD..98j3502A", "2018arXiv180206762D", "2019ApJS..242...29S", "2019JCAP...02..058G", "2019JCAP...03..008B", "2019JCAP...10..004D", "2019MNRAS.482..578B", "2019MNRAS.488.2079B", "2019MNRAS.490.4688R", "2019PhRvD.100b3516J", "2019PhRvD.100l3528J", "2020ApJ...889...89C", "2020ApJ...905..127D", "2020JCAP...05..043M", "2020JCAP...06..035J", "2020JCAP...08..007D", "2020JCAP...10..007C", "2020MNRAS.498.3403B", "2020PhRvD.101d3519H", "2020PhRvD.101h3510L", "2020PhRvD.101l3520P", "2020PhRvD.102d3516P", "2020PhRvD.102j3506A", "2020PhRvD.102l3546J", "2021JCAP...03..105B", "2021JCAP...05..069V", "2021JCAP...06..055J", "2021JCAP...11..061T", "2021MNRAS.503.2300P", "2021MNRAS.506.2780H", "2021PhRvD.103j3503S", "2021PhRvD.103j3522J", "2021PhRvD.104l3520D", "2021PhRvD.104l3529P", "2022ApJ...929....5Z", "2022ApJ...934..112W", "2022JCAP...01..033B", "2022JCAP...10..022M", "2022MNRAS.516.3029G", "2022MNRAS.tmp.2216B", "2022PhRvD.106d3515H", "2022PhRvD.106h3504T", "2022PhRvD.106j3534S", "2022arXiv220202330S", "2022arXiv221211940B", "2023JCAP...07..040G", "2023JCAP...10..028H", "2023MNRAS.525.4854B", "2023arXiv230503070G", "2023arXiv231217321W", "2024ApJ...962...21Z", "2024JCAP...05..108G", "2024MNRAS.529.2734C", "2024PhRvD.109f3504T", "2024arXiv240407249P" ]	[ "astronomy" ]	14	[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1974A&A....32..391P", "1996MNRAS.282..347M", "1997ApJ...490..493N", "1999MNRAS.308..119S", "2001ApJ...546..652S", "2001MNRAS.321..559B", "2001MNRAS.325.1312S", "2002MNRAS.335..432V", "2002PhR...367....1B", "2002PhR...372....1C", "2002PhRvD..66j3511L", "2003MNRAS.341.1311S", "2005ApJ...620..559J", "2005ApJ...634..728S", "2006ApJ...651..619J", "2006MNRAS.373..369C", "2007ApJ...664..660E", "2007PASJ...59...93N", "2010CQGra..27l4010K", "2011JCAP...10..031B", "2012JCAP...02..047G", "2012JCAP...04..019D", "2012JCAP...08..036F", "2013JCAP...12..025C", "2013JCAP...12..030C", "2013PhRvD..87l3504T", "2013PhRvD..88b3515S", "2013PhRvD..88h3502P", "2014ApJ...780..111H", "2014MPLA...2950152K", "2014PhRvD..89h3519L", "2014PhRvD..89l3522V" ]	[ "10.1088/1475-7516/2014/05/048", "10.48550/arXiv.1403.3411" ]	1403	1403.3411_arXiv.txt	Suppose that we measure a two-point correlation function (power spectrum) of density fluctuations in the Universe. We normally measure this quantity from the entire survey volume in which we have measurements of the matter distribution. Let us divide the survey volume into many subvolumes and measure the power spectrum in each subvolume. In this paper, we show that the power spectrum in each subvolume depends on environment, and is specifically correlated with the mean overdensity of that subvolume. This correlation measures how the small-scale power spectrum responds to the presence of a large-scale density fluctuation, which can be equivalently described by a non-vanishing three-point function (bispectrum). Even if the initial density fluctuations generated by inflation are perfectly Gaussian, the subsequent non-linear gravitational evolution of matter generates a non-zero bispectrum (see \cite{bernardeau/etal:2001} for a review). The ``position-dependent power spectrum'' thus offers a test of our understanding of structure formation in the Universe. Moreover, improving our understanding of structure formation increases the sensitivity to a small bispectrum generated by inflation, making it possible to test the physics of inflation using observations of the large-scale structure of the Universe. Not only is this new observable of the large-scale structure of the Universe conceptually straightforward to interpret, but it is also simpler to measure than the full bispectrum. Constraining the physics of inflation using the squeezed-limit bispectrum of the cosmic microwave background is a solved problem \cite{komatsu:2010}. However, doing the same using the bispectrum of the large-scale structure (e.g., distribution of galaxies) is considerably more challenging due to complex survey selection functions as well as to mode couplings caused by the non-linearity of the matter density field as well as the complexity of galaxy formation. This explains why only few measurements of the bispectrum have been reported in the literature \cite{scoccimarro/etal:2000,verde/etal:2001,nishimichi/etal:2006}, and further motivates our use of the position-dependent power spectrum as a simpler route to measuring the squeezed-limit bispectrum. While we mostly have galaxy redshift surveys in mind, this idea can also be applied to the projected matter density as measured through lensing. In this paper, we study this new observable. We show that the position-dependent power spectrum measures an integral of the bispectrum, which is dominated by the bispectrum in the so-called ``squeezed configurations,'' in which one wavenumber, say $k_3$, is much smaller than the other two, i.e., $k_3\ll k_1\approx k_2$. This limit of the bispectrum has a straightforward interpretation (i.e., the large-scale density fluctuation modulating the small-scale power spectrum), which can be predicted using a simple calculation. We restrict ourselves to the position-dependent power spectrum of collisionless particles in real space in this paper. We shall incorporate the effects of halo bias and redshift-space distortions in future publications. The rest of the paper is organized as follows. In \refsec{methodology}, we derive the relation between the position-dependent power spectrum, the squeezed-limit bispectrum, and the response of small-scale correlations to large-scale overdensities. In \refsec{nbody}, we present measurements of the position-dependent power spectrum from cosmological $N$-body simulations. In \refsec{modeling}, we compare various theoretical approaches to modeling the position-dependent power spectrum with the simulations. We conclude in \refsec{conclusion}. In \refapp{tr_bi_sq}, we derive the approximation of the squeezed-limit tree-level matter bispectrum.	\label{sec:conclusion} In this paper, we have proposed a novel method to measure the squeezed-limit bispectrum. By correlating the mean density fluctuation and the position-dependent power spectrum, we obtain a measurement of a certain moment of the bispectrum (integrated bispectrum) without having to actually measure three-point correlations in the data. The integrated bispectrum is dominated by the squeezed-limit bispectrum, which is much easier to model than the full bispectrum for all configurations. This is evidenced by figures~\ref{fig:sep_uni_pert}--\ref{fig:sep_uni_nonpert}, where we show model predictions accurate to a few percent using existing techniques and without tuning any parameters. A further, key advantage of this new observable is that both the mean density fluctuation and the power spectrum are significantly easier to measure in actual surveys than the bispectrum in terms of survey selection functions. In particular, the procedures developed for power spectrum estimation can be directly applied to the measurement of the position-dependent power spectrum. Additionally, the position-dependent power spectrum depends on only one wavenumber (at fixed size of the subvolume) rather than the three wavenumbers of the bispectrum. Consequently, the covariance matrix also becomes easier to model. We have measured the position-dependent power spectrum in 160 collisionless $N$-body simulations with Gaussian initial conditions, and have used two different approaches --- bispectrum modeling and the separate universe approach --- to model the measurements. All of the approaches work well on large scales, $k \lesssim 0.2~h\,{\rm Mpc}^{-1}$, and at high redshift. On small scales, where non-linearities become important, the separate universe approach (\refsec{sep_uni}) applied through the Coyote emulator prescription performs best at redshifts $z < 2$, while the SPT 1-loop predictions perform equally well at $z\geq 2$. Both show agreement to within a few percent up to $k = 0.4~h\,{\rm Mpc}^{-1}$. Accurate predictions for the position-dependent power spectrum on these and even smaller scales can be obtained by applying the separate universe approach to dedicated small-box $N$-body simulations of curved cosmologies \cite{li/hu/takada:2014}. We shall study this in an upcoming paper. The normalized integrated bispectrum is relatively insensitive to changes in cosmological parameters (\refsec{cosmodep}), and we do not expect that it will allow for competitive cosmology constraints. On the other hand, this property can also be an advantage: since this observable can be predicted accurately without requiring a precise knowledge of the cosmology, it can serve as a useful systematics test for example in weak lensing surveys. As an example, consider \refeq{iB_sq} applied to shear measurements. A constant multiplicative bias $1+m$ in the shear estimation contributes a factor $(1+m)^3$ on the left hand side of the equation, and a factor $(1+m)^4$ on the right hand side. Thus, by comparing the measured normalized integrated bispectrum with the (essentially cosmology-independent) expectation, one can constrain the multiplicative shear bias. The position-dependent power spectrum can also naturally be applied to the case of spectroscopic galaxy surveys, in which case the non-linear bias of the observed tracers also contributes to the bispectrum and position-dependent power spectrum. Thus, when applied to halos or galaxies, this observable can serve as an independent probe of the bias parameters and break degeneracies between bias and growth which are present when only considering the halo or galaxy power spectrum. We shall apply this new method to halos in $N$-body simulations, as well as to data from galaxy surveys in future papers. Finally, this approach can also be immediately applied to the projected matter density distribution as measured through weak lensing. In this case, the complexities of bias are absent and the modeling we have presented in this paper should be sufficient to describe the measurements.	14	3	1403.3411	The influence of large-scale density fluctuations on structure formation on small scales is described by the three-point correlation function (bispectrum) in the so-called ``squeezed configurations,'' in which one wavenumber, say k<SUB>3</SUB>, is much smaller than the other two, i.e., k<SUB>3</SUB> << k<SUB>1</SUB> ≈ k<SUB>2</SUB>. This bispectrum is generated by non-linear gravitational evolution and possibly also by inflationary physics. In this paper, we use this fact to show that the bispectrum in the squeezed configurations can be measured without employing three-point function estimators. Specifically, we use the ``position-dependent power spectrum,'' i.e., the power spectrum measured in smaller subvolumes of the survey (or simulation box), and correlate it with the mean overdensity of the corresponding subvolume. This correlation directly measures an integral of the bispectrum dominated by the squeezed configurations. Measuring this correlation is only slightly more complex than measuring the power spectrum itself, and sidesteps the considerable complexity of the full bispectrum estimation. We use cosmological N-body simulations of collisionless particles with Gaussian initial conditions to show that the measured correlation between the position-dependent power spectrum and the long-wavelength overdensity agrees with the theoretical expectation. The position-dependent power spectrum thus provides a new, efficient, and promising way to measure the squeezed-limit bispectrum from large-scale structure observations such as galaxy redshift surveys.	false	[ "smaller subvolumes", "k", "simulation box", "galaxy redshift surveys", "small scales", "bispectrum", "inflationary physics", "non-linear gravitational evolution", "Gaussian initial conditions", "structure formation", "lt;&lt", "the corresponding subvolume", "the mean overdensity", "the full bispectrum estimation", "collisionless particles", "large-scale structure observations", "the squeezed configurations", "large-scale density fluctuations", "the theoretical expectation", "the squeezed-limit bispectrum" ]	12.113017	2.15316	-1
477512	[ "Tokovinin, Andrei", "Mason, Brian D.", "Hartkopf, William I." ]	2014AJ....147..123T	[ "Speckle Interferometry at SOAR in 2012 and 2013" ]	66	[ "Cerro Tololo Inter-American Observatory, Casilla 603, La Serena, Chile", "U.S. Naval Observatory, 3450 Massachusetts Avenue, Washington, DC, USA", "U.S. Naval Observatory, 3450 Massachusetts Avenue, Washington, DC, USA" ]	[ "2014A&A...570A.127M", "2014A&A...570A.128E", "2014A&A...572A..91T", "2014MNRAS.443.3082T", "2015A&A...582A..25A", "2015AJ....149....8T", "2015AJ....149..195T", "2015AJ....150...50T", "2015AJ....150..130R", "2015AJ....150..136H", "2015AN....336..378A", "2015ApJ...799....4R", "2015ApJ...808...88R", "2015ApJ...809..132C", "2015ApJ...809..134P", "2015ApJ...809..135S", "2016A&A...588A..55P", "2016AJ....151...83C", "2016AJ....151..153T", "2016AJ....152..138T", "2016ApJ...831..151T", "2016SPIE.9907E..0JH", "2016SPIE.9907E..3BH", "2017AJ....153..212H", "2017AJ....154..110T", "2017AJ....154..187M", "2017AJ....154..273C", "2017ASPC..510..372H", "2017ApJ...836..139F", "2017MNRAS.469.1096D", "2017MNRAS.470.1894K", "2018AJ....155..215M", "2018AJ....155..235T", "2018ApJ...865...77R", "2018MNRAS.475.1725G", "2018MNRAS.481.5307G", "2018RAA....18...72M", "2018RMxAC..50...56M", "2019A&A...630A.119M", "2019AJ....157...91T", "2019AJ....157..249G", "2019AJ....158...48T", "2019ApJ...871...63M", "2019ApJ...885..147K", "2019MNRAS.482..471F", "2019NatAs...3..278K", "2020A&A...639A..81B", "2020AJ....159..266M", "2020AJ....160....7T", "2020gbar.conf...79M", "2021A&A...655A..15Z", "2021AIPC.2335i0002M", "2021AJ....161..155M", "2021AJ....162...41T", "2021AJ....162...53M", "2021AJ....162..156M", "2021ApJ...916..113R", "2021MNRAS.501.4903F", "2021PASA...38....2A", "2022AJ....164...58T", "2022ApJ...927L..31R", "2023AJ....165..221A", "2023AJ....166..139M", "2023AJ....166..211D", "2023RAA....23k5005M", "2024ApJS..271...55K" ]	[ "astronomy" ]	6	[ "binaries: general", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1955JO.....38...17M", "1961CiUO..120..380V", "1964ROCi..123R..59F", "1969BOBeo..27...33P", "1976MNRAS.174P..75W", "1978ApJS...37...71H", "1978ApJS...37..515H", "1978BOBeo.129....9P", "1979BAORB...9..116D", "1981ApJS...45..559H", "1983BAORB...9..262N", "1984A&AS...56....5H", "1984AJ.....89.1068H", "1984ApJ...284..806H", "1986A&AS...64....1H", "1986A&AS...65..411H", "1986A&AS...65..551B", "1988A&AS...72..543H", "1990A&AS...82...65H", "1991A&AS...90..311H", "1993A&AS...98..209H", "1994ApJS...91..793D", "1995ApJS...99..693H", "1996AJ....111..370H", "1996AJ....111..412H", "1997ASSL..223.....D", "1997ApJS..111..335H", "1999A&A...341..121S", "1999A&A...351..619F", "1999AstL...25..669T", "1999AstL...25..797B", "2000A&A...355..581H", "2000AJ....119.3084H", "2000AN....321..255S", "2001AJ....122.3466M", "2001AJ....122.3472H", "2001AJ....122.3480H", "2001SerAJ.163....5O", "2002AJ....123.1023S", "2002AJ....123.3442H", "2004A&A...415..259O", "2004A&A...418..989N", "2004A&A...425..997B", "2004AJ....127..539M", "2005A&A...433..591B", "2005A&A...441..695T", "2005AJ....129.2420M", "2005MNRAS.356.1362D", "2005RMxAA..41...17O", "2006A&A...460..695T", "2006AJ....131.1000H", "2006AJ....132.2219M", "2007AN....328..146S", "2008AJ....136..312H", "2008AJ....136..554T", "2008AJ....136.1746C", "2008ApJ...689..416B", "2008NewA...13..125C", "2008PASP..120..170T", "2009AJ....137.3358M", "2009AJ....138.1159D", "2010A&A...523A..73E", "2010AJ....139..743T", "2010AJ....139.1521L", "2010AJ....140..735M", "2010AJ....140.1623M", "2010AN....331..304C", "2010ApJS..190....1R", "2010PASP..122.1483T", "2011A&A...527A.140R", "2011AJ....141..116C", "2012AJ....143...42H", "2012AJ....144....7T", "2012AJ....144...56T", "2012AJ....144..102T", "2012JApA...33...29G", "2012JDSO....8..127R", "2012MNRAS.422.2765D", "2013AJ....145...76T", "2013AJ....146....8T" ]	[ "10.1088/0004-6256/147/5/123", "10.48550/arXiv.1403.4970" ]	1403	1403.4970_arXiv.txt	Knowledge of binary-star orbits is of fundamental value to many areas of astronomy. They provide direct measurements of stellar masses and distances, inform us on the processes of star formation through statistics of orbital elements, and allow dynamical studies of multiple stellar systems, circumstellar matter, and planets. A large fraction of visual binaries are late-type stars within 100\,pc, amenable to searches for exo-planets. However, the current orbit catalog contains some poor or wrong orbital solutions based on insufficient data. To improve the situation, we provide here new observations, revise some orbits, and compute new ones. This paper continues the series of speckle interfero\-metry observations published by \citet[][hereafter TMH10]{TMH10}, \citet{SAM09}, \citet{Hrt2012a}, and \citet{Tok2012b}. We used the same equipment and data reduction methods. All observations were obtained with the 4.1-m SOAR telescope located at Cerro Pach\'on in Chile. Our program is focused on close binaries with fast orbital motion, where the frequency of measurements (rather than the time span) is critical for orbit determination. Some of those binaries were discovered by visual observers, but most are recent discoveries made by the {\em Hipparcos} mission and by speckle interferometry, including our work at SOAR. Spectroscopic orbits are available for several fast nearby binaries resolved here. In addition, we measured close binaries with known orbits to verify and improve them when necessary, and wider pairs for calibration and quality control. Data on binary-star measures and orbits are collected in the Washington Double Star Catalog, WDS \citep{WDS}\footnote{See current version at \url{http://ad.usno.navy.mil/wds/}} and associated archives such as the {\em 4$^{\it th}$ Catalog of Interferometric Measurements of Binary Stars}, INT4 \citep{INT4}\footnote{\url{http://ad.usno.navy.mil/wds/int4.html}} and the {\it 6$^{th}$ Orbit Catalog of Orbits of Visual Binary Stars}, VB6 \citep{VB6}.\footnote{\url{http://ad.usno.navy.mil/wds/orb6.html}} These resources are extensively used here. Section~\ref{sec:data} recalls the observing technique and presents new measures, discoveries, and non-resolutions. New and updated orbits of 13 and 45 systems are given in Section~\ref{sec:orb-new} and Section~\ref{sec:orb-rev}, respectively.		14	3	1403.4970	We report the results of speckle runs at the 4.1 m Southern Astronomical Research telescope in 2012 and 2013. A total of 586 objects were observed. We give 699 measurements of 487 resolved binaries and upper detection limits for 112 unresolved stars. Eleven pairs (including one triple) were resolved for the first time. Orbital elements have been determined for the first time for 13 pairs; orbits of another 45 binaries are revised or updated. <P />Based on observations obtained at the Southern Astrophysical Research (SOAR) telescope, which is a joint project of the Ministério da Ciência, Tecnologia, e Inovação (MCTI) da República Federativa do Brasil, the U.S. National Optical Astronomy Observatory (NOAO), the University of North Carolina at Chapel Hill (UNC), and Michigan State University (MSU).	false	[ "Michigan State University", "da República Federativa", "Chapel Hill", "North Carolina", "República Federativa", "MSU", "Southern Astronomical Research", "UNC", "the U.S. National Optical Astronomy Observatory", "upper detection limits", "NOAO", "Brasil", "MCTI", "Inovação", "Tecnologia", "the Ministério da Ciência", "e Inovação (MCTI", "the University of North Carolina", "first", "orbits" ]	6.961602	11.854428	-1
437362	[ "Suad, L. A.", "Caiafa, C. F.", "Arnal, E. M.", "Cichowolski, S." ]	2014A&A...564A.116S	[ "A new catalog of H i supershell candidates in the outer part of the Galaxy" ]	25	[ "Instituto Argentino de Radioastronomía (IAR), CC 5, 1894, Villa Elisa, Argentina", "Instituto Argentino de Radioastronomía (IAR), CC 5, 1894, Villa Elisa, Argentina; Facultad de Ingeniería, Universidad de Buenos Aires, C.A.B.A., C117AAZ, Buenos Aires, Argentina", "Instituto Argentino de Radioastronomía (IAR), CC 5, 1894, Villa Elisa, Argentina; Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, 1900, La Plata, Argentina", "Instituto de Astronomía y Física del Espacio (IAFE), Ciudad Universitaria, C.A.B.A, 1428, Buenos Aires, Argentina" ]	[ "2014MNRAS.444.3657P", "2015A&ARv..23....3D", "2015AJ....149...12C", "2015AJ....149..189S", "2015ApJ...807...68J", "2016A&A...585A.154S", "2016A&A...587A...5E", "2016ApJ...827...42P", "2016BAAA...58..209S", "2017ApJ...850...71F", "2017MNRAS.465.1720Y", "2018A&A...619A.101E", "2018MNRAS.479.5620T", "2019A&A...624A..43S", "2019A&A...628A..44F", "2020AJ....160...66P", "2021A&A...652A.142P", "2022A&A...668A..44S", "2022BAAA...63..143S", "2023AJ....165...16C", "2023ApJ...944L...8S", "2023BAAA...64..105C", "2023BAAA...64..121S", "2023arXiv231212210F", "2024A&A...685L..12L" ]	[ "astronomy" ]	8	[ "ISM: bubbles", "ISM: structure", "methods: data analysis", "techniques: image processing", "radio lines: ISM", "Astrophysics - Astrophysics of Galaxies" ]	[ "1973A&AS....8....1W", "1974A&AS...14....1H", "1977ApJ...218..377W", "1979ApJ...229..533H", "1981ApJ...248..119H", "1984ApJS...55..585H", "1987ApJ...317..190M", "1988ARA&A..26..145T", "1989ApJ...342..272F", "1990ARA&A..28..215D", "1997agnh.book.....H", "1998ApJ...501L.163E", "1998ApJ...503L..35L", "1998MNRAS.298..797T", "2001ApJ...551..394M", "2002ApJ...578..176M", "2003AJ....125.3145T", "2005A&A...437..101E", "2005A&A...440..775K", "2007A&A...476..255A", "2007ApJ...661..285D", "2013A&A...550A..23E" ]	[ "10.1051/0004-6361/201323147", "10.48550/arXiv.1403.4141" ]	1403	1403.4141_arXiv.txt	When viewed in the neutral hydrogen (\hi) line emission ($\lambda = 21$ cm), the interstellar medium (ISM) reveals a complex structuring that manifests itself by the presence of a plethora of structures such as shells, supershells, filaments, arcs, cavities, worms, and loops. In particular, \hi\, shells and supershells are detected in a given radial velocity range as voids in the \hi\, emission distribution that are surrounded completely, or partially, by walls of enhanced \hi\, emission. The physical processes usually invoked to explain the formation of \hi\, shells and supershells are the combined action of strong winds from massive stars upon the surrounding interstellar medium, and their ultimate explosion as supernovae. Other alternative mechanisms have also been proposed to explain their creation, such as gamma-ray bursts \citep{efr98,loe98} and the infall of high velocity clouds \citep{ten88}. Several catalogs of shells and supershells have already been constructed with a variety of different techniques such as visual identification \citep{hei79, hei84, mcc02} and automatic identification algorithms \citep{ehl05,dai07,ehl13}. Based on a visual inspection of photographic representations of the \hi\, emission distribution derived from the \hi\, survey of \citet{wea73}, a first catalog of \hi\, shells and supershells was constructed by \citet{hei79}, who defines as supershells to those structures requiring at least $3 \times 10^{52}$ erg for their creation. A total of 63 structures were identified. Later, using the 21 cm database of \cite{hei74} for $\vert b \vert > 10^\circ$, \citet{hu81} discovered 50 new shells. In a subsequent work, \citet{hei84} combined the surveys of \cite{wea73} and \cite{hei74} in order to eliminate the boundary problems at $\vert b \vert = 10^\circ$, in this way finding a total of 42 new structures. Afterwards, \citet{mcc02} report the discovery of 19 new \hi\, shells, using the Southern Galactic Plane Survey (SGPS) \citep{mcc01b}. On the other hand, applying a fully automatic method to the \hi\, Leiden-Dwingeloo survey \citep{har97}, \cite{ehl05} searched for regions having a local minimum in the \hi\, emissivity that were completely encircled by regions of higher emissivity. Later on, the same authors \citep{ehl13} used the Leiden-Argentine-Bonn (LAB) survey \citep{kal05} to conduct an all-sky search for \hi\, shells, employing an automatic procedure that is slightly different from the one described in \cite{ehl05}. Using an artificial neural network and the \hi\, database of the Canadian Galactic Plane Survey (CGPS) \citep{tay03}, \cite{dai07} identified a large number of small expanding \hi\, shells in the Perseus arm. Bearing in mind that different identification techniques applied to the same database \citep[see][]{ehl05,ehl13} provide different results and that quite different structures are identified in the same region of the sky when different identification criteria and databases are used (see feature GSH\, 263+00+47 of \citeauthor{mcc02} \citeyear{mcc02} and GS\,263-02+45 of \citeauthor{arn07} \citeyear{arn07}), we believe that it is worth making the effort to unveil the presence of these structures (shells and supershells) by using a relatively different approach for their identification. In particular, in this paper we deal solely with the identification of the so-called \hi\, supershells located in the outer part of the Galaxy by using a novel procedure that combines both visual and automatic algorithms of identification. The constraint of only analyzing the second and third galactic quadrants stems from the fact that toward this part of the Galaxy the radial velocity-distance relationship only has a single value, a fact that strongly simplifies determination of the physical sizes of the detected structures. The paper is organized as follows. In Section 2 we briefly describe the database used in constructing the catalog of \hi\, supershell candidates, while in Section 3 the identification techniques are described in some detail. Section 4 describes the selection effects of our catalog. The statistical properties of the supershell candidates are presented in Section 5, and a comparison with previous catalogs is made in Section 6.	A new catalog of \hi\, supershell candidates has been constructed using a combination of an automatic detection algorithm plus a visual one. It is known that pure visual identification methods are difficult and very time consuming mainly because the dynamic range has to be adjusted quite often to make structures visible. However, the eye is an incredibly powerful instrument, especially when images are irregular, since it is able to combine disconnected patterns. Our algorithm was trained on an initial visual catalog and then applied to the whole dataset. At the end of this process, all the detected structures were again carefully inspected by eye to do the final selection. In other words, in our method, we have combined the power of visual inspection, together with the power of a computer-based algorithm working in a supervised (semi-automatic) mode to optimize the time required to analyze a huge amount of data. A total of 566 candidate structures has been detected, 347 in the second Galactic quadrant and 219 in the third one. From the distribution of the detected structures in the sky, it can be seen that for the second Galactic quadrant, we have detected structures with distances up to 32 kpc form the Sun, while for the third quadrant all the structures have distances less than 17 kpc from the Sun. The explanation of this effect is far from clear. It shows that the outer part of the Galaxy deserves a thorough study. Owing to our selection criteria, the catalog suffers from selection effects. Bearing this in mind only features closer to the Sun than 5.7 kpc were used to derive their statistical properties. The estimated mean weighted effective radius is about 160 pc for both Galactic quadrants. The derived eccentricities indicate that about $98 \, \%$ of the supershells are elliptical. The mean weighted eccentricity is 0.8 $\pm$ 0.1 for both Galactic quadrants. An inspection of the orientation angle values (e.g., the angle between the Galactic plane and the feature's major semi-axis) shows that most of the supershell candidates are elongated parallel to the Galactic plane, in agreement with the conclusions of other researchers. Based on this finding, it is believed that the galactic density gradient plays a minor, if any, role in the time evolution of these structures. On the other hand, the magnetic field running parallel to the Galactic disk could be the responsible for the observed effect. Otherwise, the major axes of the structures should be predominantly perpendicular to the Galactic plane. The distribution of the centroid of the structures relative to the Galactic plane shows that roughly $\sim$ 40\% of the supershells are confined to $z \leq 500 $ pc. After applying correction factors owing to the incompleteness of our sample, we derived the surface density of \hi\, supershells. Though it decreases as the galactocentric distance increases, in close agreement with the findings of other researchers, the actual surface density in the solar neighborhood is almost a factor of 2 higher, 7.4 kpc$^{-2}$, than the one derived before by \cite{ehl13}. This could be a direct consequence of the ``ability'' of our method to identify incomplete features. In line with this, we would like to point out that roughly only half of the supershell candidates are found to be completely closed. The decrease in the surface density with galactocentric distance is well fit by a Gaussian function with a radial scale length of 4.4 $\pm$ 0.3 kpc. The surface (f$_{2D}$) and volume (f$_{3D}$) filling factors are f$_{2D} = 0.5 \pm 0.1$ and f$_{3D} = 0.04^{+0.01}_{-0.02}$. As mentioned above, because our algorithm is able to detect incomplete structures, we have, as a byproduct, a catalog of ``galactic chimney'' candidates, which contains 80 structures. A clear relationship between the effective radius of the structures and their expansion velocities is not detected in our catalog. Under the assumption that their genesis is mostly consequence of the action of massive stars (stellar winds and supernova explosions), the dynamic age of the structures are between 5-50 Myr and most of the structures fall between values of $L_w/n$ of (0.1-100) $\times 10^{36}$ erg/s cm$^{-3}$. A comparison with the structures listed in our catalog with those given in other shell/supershell catalogs, shows that we identify, on average, 57\% of the structures listed elsewhere. The lowest correspondence is with the catalog of \cite{hei79} ($\sim$ 36\%), and the highest is with \cite{mcc02} ($\sim$ 75 \%).	14	3	1403.4141	<BR /> Aims: The main goal of this work is to a have a new neutral hydrogen (H i) supershell candidate catalog to analyze their spatial distribution in the Galaxy and to carry out a statistical study of their main properties. <BR /> Methods: This catalog was carried out making use of the Leiden-Argentine-Bonn (LAB) survey. The supershell candidates were identified using a combination of two techniques: a visual inspection plus an automatic searching algorithm. Our automatic algorithm is able to detect both closed and open structures. <BR /> Results: A total of 566 supershell candidates were identified. Most of them (347) are located in the second Galactic quadrant, while 219 were found in the third one. About 98% of a subset of 190 structures (used to derive the statistical properties of the supershell candidates) are elliptical with a mean weighted eccentricity of 0.8 ± 0.1, and ~70% have their major axes parallel to the Galactic plane. The weighted mean value of the effective radius of the structures is ~160 pc. Owing to the ability of our automatic algorithm to detect open structures, we have also identified some "galactic chimney" candidates. We find an asymmetry between the second and third Galactic quadrants in the sense that in the second one we detect structures as far as 32 kpc, while for the 3rd one the farthest structure is detected at 17 kpc. The supershell surface density in the solar neighborhood is ~8 kpc<SUP>-2</SUP>, and decreases as we move farther away form the Galactic center. We have also compared our catalog with those by other authors. <P />Full table is only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/564/A116">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/564/A116</A>	false	[ "open structures", "structures", "candidate catalog", "~8 kpc", "Galactic", "the farthest structure", "H i", "second", "anonymous ftp", "pc", "third", "3rd", "566 supershell candidates", "The supershell candidates", "the supershell candidates", "the second Galactic quadrant", "both closed and open structures", "190 structures", "the structures", "%" ]	11.17429	9.174909	-1
483253	[ "Bisterzo, S.", "Travaglio, C.", "Gallino, R.", "Wiescher, M.", "Käppeler, F." ]	2014ApJ...787...10B	[ "Galactic Chemical Evolution and Solar s-process Abundances: Dependence on the <SUP>13</SUP>C-pocket Structure" ]	221	[ "INAF-Astrophysical Observatory Turin, Turin, Italy; Department of Physics, University of Turin, Turin, Italy;", "INAF-Astrophysical Observatory Turin, Turin, Italy; B2FH Association-c/o Strada Osservatorio 20, I-10023 Turin, Italy.;", "Department of Physics, University of Turin, Turin, Italy; B2FH Association-c/o Strada Osservatorio 20, I-10023 Turin, Italy.", "Joint Institute for Nuclear Astrophysics (JINA), Department of Physics, University of Notre Dame, Notre Dame IN, USA", "Karlsruhe Institute of Technology, Campus Nord, Institut für Kernphysik, Karlsruhe, Germany" ]	[ "2014A&A...565A.123M", "2014A&A...568A..47H", "2014A&A...569A..43M", "2014A&A...570A..76B", "2014A&A...572A.108T", "2014ApJ...787...41T", "2014ApJ...788..163L", "2014ApJ...795..141T", "2014ApJ...797...44F", "2014ApJ...797..123H", "2014MNRAS.443L..99G", "2015A&A...583A..67J", "2015A&A...584A..86S", "2015AcASn..56..564Z", "2015ApJ...801...53C", "2015ApJ...803...12L", "2015ApJ...808L..10C", "2015ApJ...809..180M", "2015EL....10962001W", "2015EPJWC..9302013R", "2015GeCoA.165..361B", "2015MNRAS.446.3651M", "2015MNRAS.449..506B", "2015MNRAS.452.1970W", "2015MNRAS.452.2804S", "2015RAA....15.1264L", "2016A&A...586A...1V", "2016A&A...586A..49B", "2016A&A...589A..17Y", "2016A&A...591A..53B", "2016A&A...592A..29M", "2016A&A...593A..79S", "2016A&A...593A.125S", "2016ApJ...818..125T", "2016ApJ...819..118T", "2016ApJ...825...26K", "2016ApJ...827...29M", "2016ApJ...833..225Z", "2016ApJS..225...24P", "2016ChA&A..40..322Z", "2016JPhCS.728b2002L", "2016LPI....47.1888N", "2016MNRAS.459.4174G", "2016MmSAI..87..229K", "2016Natur.537..399B", "2016PASA...33...28B", "2016PASA...33...40M", "2016PhRvC..93e5803T", "2016SciA....2E1400T", "2017A&A...598A..54H", "2017A&A...601A..56T", "2017A&A...606A..55M", "2017A&A...606A..94D", "2017A&A...607A..91H", "2017A&A...608A..89M", "2017A&A...608A.145B", "2017AcASn..58...47Z", "2017AcASn..58...61Z", "2017ApJ...835...97B", "2017ApJ...846...23O", "2017ApJ...850...80R", "2017ApJ...850L..12T", "2017E&PSL.473..215P", "2017EPJD...71..169R", "2017GeocJ..51...17F", "2017IAUS..323...86L", "2017MNRAS.466.4672F", "2017MNRAS.469.4378M", "2017PhRvC..95c5803L", "2018A&A...611A..30S", "2018A&A...615A.172M", "2018A&A...617A.106M", "2018A&A...619A.143G", "2018ARA&A..56..223B", "2018ApJ...854..105C", "2018ApJ...855...83H", "2018ApJ...856..138J", "2018ApJ...863..115V", "2018ApJ...865..112L", "2018ApJS..236...36R", "2018ChA&A..42..538Z", "2018EPJA...54..221L", "2018EPJWC.16501039P", "2018GeCoA.221...21P", "2018GeCoA.221...37P", "2018GeCoA.239..346E", "2018IAUS..330..331M", "2018MNRAS.474.2580S", "2018MNRAS.476..690P", "2018MNRAS.476.3432P", "2018OAP....31...78M", "2019A&A...622A..74J", "2019A&A...626A..15H", "2019A&A...631A.113F", "2019A&A...631A.171S", "2019AIPC.2076c0002K", "2019ApJ...875..106C", "2019ApJ...877..101S", "2019ApJ...882...40J", "2019ApJ...883...62Y", "2019ApJ...885..128O", "2019ApJ...886...84G", "2019AstL...45..341M", "2019GeCoA.261..145B", "2019GeCoA.265..413C", "2019M&PS...54.1092B", "2019MNRAS.482.1690B", "2019MNRAS.484.3846M", "2019MNRAS.485.3623R", "2019MNRAS.486.5349D", "2019MNRAS.487.1745W", "2019MNRAS.489.1697M", "2019MNRAS.490.2219N", "2019Natur.574..497W", "2019PrPNP.107..109K", "2019Sci...363..474J", "2020A&A...633L...9J", "2020A&A...634A..72S", "2020A&A...634A..84S", "2020A&A...634L...2S", "2020A&A...635A.104H", "2020A&A...638A..64P", "2020A&A...641A.127R", "2020A&A...642A.227A", "2020A&A...644A..19W", "2020ApJ...889..119T", "2020ApJ...893...37R", "2020ApJ...896...64L", "2020ApJ...899..133H", "2020ApJ...902L..24R", "2020EPJWC.23907003M", "2020GeCoA.270..475E", "2020MNRAS.491.1832P", "2020PhLB..80435405M", "2020PhRvC.101d5802D", "2020arXiv200604629M", "2020arXiv201004632W", 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"2023A&A...673A.155Q", "2023A&A...675A.194S", "2023A&A...677A..22R", "2023A&A...677A..60C", "2023A&A...677A..61M", "2023A&A...677A..74D", "2023ApJ...944..121W", "2023ApJ...944..123V", "2023ApJ...952L..25S", "2023ApJ...954...55R", "2023ApJ...956....7M", "2023ApJ...957...10A", "2023EPJA...59..302P", "2023EPJP..138..709B", "2023GeCoA.344...40B", "2023JQSRT.29408392A", "2023MNRAS.518.5414R", "2023MNRAS.523.2111M", "2023MNRAS.523.2974M", "2023MNRAS.525.1329L", "2023PhRvC.107b5805G", "2023Sci...382.1177R", "2023arXiv230209188F", "2023arXiv231202258S", "2024A&A...683A.200S", "2024A&A...684A..15N", "2024AJ....167...98G", "2024AJ....167..184R", "2024ApJ...962...33V", "2024ApJ...964...96N", "2024ApJ...966...11P", "2024ApJ...966L..37H", "2024MNRAS.527.1389H", "2024MNRAS.528.4354R", "2024MNRAS.530.2512L", "2024PhLB..85138581S", "2024arXiv240500096S", "2024arXiv240500110C", "2024arXiv240513641E", "2024arXiv240519330V", "2024arXiv240602691R" ]	[ "astronomy" ]	13	[ "stars: AGB and post-AGB", "Galaxy: evolution", "Sun: abundances", "Astrophysics - Solar and Stellar Astrophysics" ]	[ "1957AJ.....62....9C", "1957RvMP...29..547B", "1965ApJS...11..121S", "1967ApJ...148...69C", "1983ARA&A..21..271I", "1988ApJ...328..653B", "1989GeCoA..53..197A", "1990ApJ...348..357T", "1990ApJS...72..387S", "1991ApJ...373L...5H", "1992ApJ...387..263R", "1994A&A...283..593P", "1994ApJ...437..396K", "1995ApJ...440L..85S", "1995ApJ...450..302L", "1997A&A...324L..81H", "1998ApJ...497..388G", "1999A&A...346L..37L", "1999ARA&A..37..239B", "1999ApJ...521..691T", "1999ApJ...523..234T", "1999ApJ...524..226D", "1999ApJ...525..886A", "2000A&A...362..599G", "2000ADNDT..76...70B", "2000MmSAI..71..719S", "2001ApJ...549..346T", "2001ApJ...551..833I", "2001ApJ...557..802B", "2001ApJ...559.1117A", "2001Natur.412..793V", "2001PhRvC..64a5801B", "2001PhRvL..87t2501J", "2002A&A...390..235N", "2002ApJ...576..323R", "2002ApJ...579..817A", "2002ApJ...580.1149A", "2002PhRvC..66f4603R", "2003ApJ...582.1251R", "2003ApJ...591..936S", "2003ApJ...593..486L", "2003ApJ...593.1056H", "2003ApJS..148..175S", "2003MNRAS.340..722D", "2003PASA...20..279K", "2003PASA...20..389S", "2004A&A...415.1089S", "2004A&A...421L..25G", "2004A&A...425.1029T", "2004ARA&A..42...39C", "2004ApJ...601..864T", "2004ApJ...605..425H", "2004PhRvC..69b5803P", "2004PhRvC..69e5801W", "2005A&A...430..655S", "2005A&A...431..279V", "2005A&A...439.1075V", "2005AIPC..769.1327M", "2005ARA&A..43..435H", "2005ARA&A..43..481A", "2006A&A...451..651J", "2006A&A...454..917A", "2006ApJ...642.1082M", "2006ApJ...647..685W", "2006NuPhA.777..311S", "2007A&A...461..641R", "2007ApJ...659.1265S", "2007ApJ...671..802B", "2008ARA&A..46..241S", "2008ApJ...673..434H", "2008ApJ...673..854Y", "2008ApJ...681.1524L", "2008ApJ...685.1089H", "2008ApJS..174..158S", "2009A&A...508..909Z", "2009ApJ...696..797C", "2009ApJ...705L..31G", "2009MNRAS.396.1046L", "2009PASA...26..153S", "2009PASA...26..243H", "2010A&A...512A..10S", "2010ApJ...710.1557P", "2010MNRAS.403.1413K", "2010MNRAS.404.1529B", "2010PhRvC..81e5801T", "2010PhRvC..82a5802M", "2010PhRvC..82c5806M", "2011ApJ...739...93T", "2011ApJS..197...17C", "2011MNRAS.418..284B", "2011PhRvC..84a5801T", "2011PhRvC..84e5802T", "2011PrPNP..66..346T", "2011RvMP...83..157K", "2012A&A...538L...2F", "2012A&A...540A..44V", "2012A&A...545A..31H", "2012ApJ...747....2L", "2012ApJ...747...53M", "2012ApJ...751....8K", "2012ApJ...755L..10R", "2012ApJ...756...36R", "2012ApJ...757..164R", "2012PhRvC..85c5802R", "2012PhRvC..85d4615M", "2012PhRvC..85f5809L", "2013A&A...551A..57H", "2013A&A...552A.128M", "2013A&A...555L...3G", "2013ApJ...762...31P", "2013ApJ...764...78R", "2013ApJ...774...98P", "2013ApJ...776...59D", "2013ApJS..208...20B", "2013JPhG...40a3201A", "2013MNRAS.431.3642V", "2013PhRvC..88c5802K", "2013PhRvL.110b2501L", "2013arXiv1311.5381B", "2014ApJ...780...95L", "2014MNRAS.437..195D" ]	[ "10.1088/0004-637X/787/1/10", "10.48550/arXiv.1403.1764" ]	1403	1403.1764_arXiv.txt	The solar-system abundances result from contributions of different nucleosynthesis processes. Isotopes heavier than the iron group are produced via neutron captures, the $slow$ and the $rapid$ processes ($s$ and $r$). Exceptions include the 35 or so\footnote{The original 35 proton-rich nuclei reported by \citet{cameron57} and \citet{burbidge57} were defined on the basis of being apparently excluded by the $s$- or by the $r$-process.} isotopes synthesized by the $p$-process. \\ The origin of the $r$-process is currently attributed to explosive nucleosynthesis in massive stars, even if the astrophysical conditions are still not well defined (see, e.g., \citealt{thielemann11}). The $p$-process is identified with photo-disintegration processes (the so-called $\gamma$-process; \citealt{howard91}) of heavy elements occurring in core collapse Supernovae of Type II (e.g., \citealt{rauscher02,hayakawa08}). Accreting white dwarfs in binary systems with masses approaching the Chandrasekhar limit may substantially contribute to the production of $p$-process nuclei when they explode as Type Ia Supernovae (e.g., \citealt{travaglio11}, and references therein). \\ Concerning the $s$-process, pioneering studies \citep{seeger65,clayton67} have demonstrated that the solar $s$-distribution of isotopes from iron to bismuth may be reproduced by three components: the $weak$, the $main$ and the $strong$ component. The $weak$ component contributes to $s$-isotopes up to A $\sim$ 90. It takes place in massive stars during core He and convective shell C burning (see, e.g., \citealt{pignatari10}; \citealt*{fris12}). The $main$ component (A $\sim$ 90 to 208) derives from low mass asymptotic giant branch (AGB) stars, during their thermally pulsing (TP) phase \citep{arlandini99,goriely00,cristallo09,bisterzo11,lugaro12}. The $strong$ component is needed to reproduce about half of solar $^{208}$Pb. It originates from AGB stars of low metallicity (at [Fe/H] $\sim$ $-$1; \citealt{gallino98,travaglio01,vaneck01}). After a limited number of TPs, at the quenching of recurrent thermal instabilities in the He shell, the convective envelope penetrates in the outer region of the He-intershell (third-dredge-up, TDU), bringing to the surface newly synthesized $^{12}$C and $s$-process isotopes. The whole envelope undergoes strong mass loss by stellar winds, leaving the degenerate core that eventually will become a white dwarf. \\ The major neutron source of the $s$-process in low mass AGB stars ($M$ $<$ 4 $M_\odot$) is the $^{13}$C($\alpha$, n)$^{16}$O reaction, which is activated in radiative conditions during the interpulse phase \citep{straniero95}. At the quenching of a TDU, a small amount of hydrogen is assumed to penetrate from the envelope into the top layers of the radiative He-rich zone. At hydrogen re-ignition, a thin $^{13}$C-pocket forms by proton captures on the abundant $^{12}$C. When the temperature of $\sim$9$\times$10$^7$ K is reached (which corresponds to $kT$ $\sim$ 8 keV), $^{13}$C is fully depleted releasing neutrons with a density of about 10$^7$ n/cm$^{-3}$. \\ The $^{13}$C in the pocket is of primary origin, directly synthesized in the star from the $^{12}$C produced by partial He burning during previous thermal pulse and is independent of the initial composition. However, the number of free neutrons per iron seed increases with decreasing metallicity. Consequently, for a given $^{13}$C-pocket strength, the $s$-process flow firstly feeds the $s$-process peak Sr-Y-Zr (at magic neutron number N = 50), extending up to $^{136}$Ba, then reaches the second $s$-process peak (Ba-La-Pr-Ce-Nd at N = 82), extending up to $^{204}$Pb-$^{207}$Pb, with a progressive increasing ratio of the heavy-$s$ elements (hs) to the light-$s$ elements (ls). At even lower metallicity, it mainly feeds $^{208}$Pb (N = 126) at the termination of the $s$-process path. Therefore, the $s$-process in AGB stars is strongly metallicity dependent. The complexity of the $s$-process distribution in AGB stars is confirmed by spectroscopic observations in different stellar populations (planetary nebulae, post-AGB, MS, S, C(N), Ba, CH and CEMP-$s$ stars; see, e.g., \citealt{smith90,pequignot94,aoki02,abia02,allen06,jonsell06,sharpee07,sterling08,reyniers07,zamora09}). Furthermore, for a given metallicity a spread in the $s$-process distribution is observed for each class of stars. A range of $^{13}$C-pocket strengths is needed in order to explain this spread (see, e.g., \citealt{busso01,sneden08,kaeppeler11,bisterzo11,lugaro12}). A similar spread is shown by $s$-process isotopic signatures found in presolar meteoritic SiC grains, which originated in the outflow of AGB stars (see, e.g., \citealt{lugaro03,clayton04,zinner07}). A marginal activation of the $^{22}$Ne($\alpha$, n)$^{25}$Mg reaction occurs at the bottom of the advanced convective thermal pulses, where a temperature of $T$ $\sim$3$\times$10$^8$ K is reached. A short neutron burst is released with a peaked neutron density (up to $N_n$$\sim$10$^{11}$ n/cm$^{-3}$), which provides only a few percent of the total neutron exposure, but affects the abundances of some important isotopes close to the main branchings of the $s$-process (e.g., $^{85}$Kr and $^{95}$Zr, sensitive to neutron density). For an extended review on the $s$-process we refer to \citet{kaeppeler11}, and references therein. Several model uncertainties affect the AGB yields, e.g., the mass loss efficiency, the deepness of the TDU, and the number of TPs (\citealt{herwig05,straniero06}). Particularly challenging in AGB modeling is the formation of the $^{13}$C-pocket and the physical prescriptions involved. \citet{iben83} suggested that, as a consequence of the TDU, a sharp discontinuity between the H-rich envelope and the He- and C-rich intershell forms, and a partial mixing of protons from the envelope into the He-intershell may take place. The amount of protons that diffuses into the He-intershell must be small to allow the production of $^{13}$C and to limit the further conversion of $^{13}$C to $^{14}$N by proton captures. $^{14}$N mainly acts as a neutron poison of the $s$-process via the $^{14}$N(n, p)$^{14}$C reaction. In the external zone of the pocket (where protons are more abundant), a $^{14}$N-rich zone may also form, depending on the proton profile adopted in the AGB model; this region plays a minor role in the $s$-process nucleosynthesis. \\ The hydrogen profile and, correspondingly, the internal structure of the $^{13}$C-pocket may depend on the initial mass and metallicity of the AGB, and on the physical mechanisms that may compete inside the star itself. \citet{herwig97} proposed an exponential diffusive overshooting at the borders of all convective zones. FRUITY models by \citet{cristallo09,cristallo11} adopted an opacity-induced overshooting at the base of the convective envelope by introducing in the model an exponentially decaying profile of the convective velocity. \\ Starting from \citet{langer99}, \citet{herwig03} and \citet{siess04}, rotation was introduced in stellar models to study its impact on the $^{13}$C-pocket structure. First studies agree that rotation-induced instabilities reduce the total mass of $^{13}$C in the pocket, owing to the $^{14}$N contamination in the $^{13}$C-rich layer, compromising the $s$-process efficiency. \citet{piersanti13} confirm that rotation-induced instabilities modify the mass extension of both $^{13}$C and $^{14}$N abundances in the pocket, and their overlap as well. Moreover, they suggest that meridional (Eddington-Sweet) circulation may smooth and enlarge the $^{13}$C-rich zone of the pocket. \\ \citet{denissenkov03} demonstrated that a weak turbulence induced by gravity waves presents additional alternative for the $^{13}$C-pocket formation. On the other hand, \citet{busso12} revisit the idea that rotation favors mixing indirectly, through the maintenance of magnetic dynamo mechanisms, producing the buoyancy of toroidal magnetic structures \citep{busso07}. \\ Further investigations on rotation, magnetic fields, gravity waves and the interplay between these several mechanisms will help to shed light on this challenging issue. The solar system $s$-process distribution is the result of the nucleosynthesis of all previous generations of AGB stars that have polluted the interstellar medium. Therefore, a Galactic chemical evolution (GCE) model is required to follow the complex evolutionary processes of the Milky Way. \citet{travaglio99} showed that AGB yields computed within a weighted average over the range of $^{13}$C-pockets are needed in the framework of GCE model to reproduce the $s$-distribution observed in the solar system. This is also suggested by the spectroscopic $s$-process spread observed in individual stars and the isotopic anomalies measured in presolar SiC grains. \\ GCE calculations succeeded to reproduce the solar abundances of $s$-only isotopes between $^{134}$Ba and $^{208}$Pb (\citealt{travaglio01,travaglio04}). However, a deficit (of about 25\%) between GCE predictions at the solar epoch and the abundances measured in the solar system was found for Sr, Y, Zr, and $s$-process isotopes up to A = 130, including ten $s$-only isotopes from $^{96}$Mo to $^{130}$Xe (see also \citealt{kaeppeler11}; their Fig.~15). The $weak-s$ process produces isotopes up to Sr ($\sim$10\% to Sr and $\la$5\% to Y and Zr isotopes), with a negligible contribution afterwards. An additional $r$-fraction of $\sim$10\% was estimated for solar Sr-Y-Zr, evaluating the $r$-contribution from the prototypical r-II star CS 22892--052\footnote{CS 22892--052 has an $r$-process enrichment of $\sim$40 times the solar-scaled composition ([Eu/Fe] = 1.6; \citealt{sneden03}).}. Summing up, the $s$-, $r$- (and $p$-) contributions predicted by current stellar models are not sufficient to explain the solar abundances of light isotopes from $A$ $\sim$ 90 to 130. \citet{travaglio04} hypothesized the existence of an additional process of unknown origin, called by the authors LEPP (light-element primary process), which must supply $\sim$8\% of solar Sr and $\sim$18\% of solar Y and Zr. Several scenarios have been recently explored, involving a primary component in massive stars that comes from the activation of $^{13}$C($\alpha$, n)$^{16}$O in the C core when the temperature is low enough to prevent the $^{13}$N($\gamma$, p)$^{12}$C reaction from becoming efficient (defined "cold" C-burning component or $cs$-component by \citealt{pignatari13}), and/or a light $r$-process induced by explosive stellar nucleosynthesis, e.g., in the neutrino-driven winds (\citealt{arcones13} and references therein). The aim of this work is to investigate the influence of one of the major AGB yield uncertainties, the formation of the $^{13}$C-pocket, on the predicted solar $s$-process distribution, from light neutron capture isotopes (to verify the need of LEPP) up to Pb and Bi at the end of the $s$-path. Firstly, in Section~\ref{update} we present updated $s$-percentages for isotopes from Kr to Bi with respect to \citet{travaglio04}. In Section~\ref{pocket}, we test the effect of AGB yields computed with different choices of the internal structure of the $^{13}$C-pocket on GCE $s$-predictions at the epoch of the solar-system formation. Our results are briefly summarized in Section~\ref{conclusions}.	We study the solar abundances of heavy $s$-isotopes at the epoch of the formation of the solar system as the outcome of nucleosynthesis occurring in AGB stars of various masses and metallicities. At present, one of the major uncertainties of AGB stellar model is the formation of the $^{13}$C-pocket. Our aim is to investigate the impact of uncertainties concerning the internal structure of the $^{13}$C-pocket on the GCE $s$-distribution, by carrying out different tests in which we modify the $^{13}$C and $^{14}$N abundances in the pocket, and the size in mass of the pocket itself. Thereby, we obtain that GCE $s$-process predictions at the epoch of the solar-system formation marginally depend on the choice of the $^{13}$C profile and on the mass of the pocket when a range of $^{13}$C-pocket strengths is adopted. The GCE model may reproduce within the solar error bars the $s$-contribution to isotopes with A $>$ 130. The missing contribution to isotopes in the range between A = 90 to 130 found by \citet{travaglio04} is confirmed by the present analysis: an additional $s$-process (LEPP-$s$) is required to account for the missing component of ten $s$-only isotopes ($^{96}$Mo, $^{100}$Ru, $^{104}$Pd, $^{110}$Cd, $^{116}$Sn, $^{122,123,124}$Te and $^{128,130}$Xe). Based on the tests made in this paper, the LEPP-$s$ contribution remains essentially the same, independently of the internal structure of the $^{13}$C-pocket. First indications in favor of this process have been analyzed and discussed by \citet{pignatari13}. \\ An additional primary contribution is being explored to account for a complementary light-$r$ contributions. In spite of promising theoretical improvements related to the explosive phases of massive stars and core collapse Supernovae (\citealt{arcones13} and references therein), as well as recent spectroscopic investigations \citep{roederer12,hansen12,hansen13}, a full understanding of the origin of the neutron capture elements from Sr up to Xe is still lacking.	14	3	1403.1764	We study the s-process abundances (A >~ 90) at the epoch of the solar system formation. Asymptotic giant branch yields are computed with an updated neutron capture network and updated initial solar abundances. We confirm our previous results obtained with a Galactic chemical evolution (GCE) model: (1) as suggested by the s-process spread observed in disk stars and in presolar meteoritic SiC grains, a weighted average of s-process strengths is needed to reproduce the solar s distribution of isotopes with A > 130; and (2) an additional contribution (of about 25%) is required in order to represent the solar s-process abundances of isotopes from A = 90 to 130. Furthermore, we investigate the effect of different internal structures of the <SUP>13</SUP>C pocket, which may affect the efficiency of the <SUP>13</SUP>C(α, n)<SUP>16</SUP>O reaction, the major neutron source of the s process. First, keeping the same <SUP>13</SUP>C profile adopted so far, we modify by a factor of two the mass involved in the pocket; second, we assume a flat <SUP>13</SUP>C profile in the pocket, and we test again the effects of the variation of the mass of the pocket. We find that GCE s predictions at the epoch of the solar system formation marginally depend on the size and shape of the <SUP>13</SUP>C pocket once a different weighted range of <SUP>13</SUP>C-pocket strengths is assumed. We obtain that, independently of the internal structure of the <SUP>13</SUP>C pocket, the missing solar system s-process contribution in the range from A = 90 to 130 remains essentially the same.	false	[ "updated initial solar abundances", "the missing solar system s-process contribution", "the s-process abundances", "s-process strengths", "presolar meteoritic SiC grains", "the solar system formation", "the solar s-process abundances", "the s process", "isotopes", "the s-process spread", "different internal structures", "<SUP>13</SUP>C-pocket strengths", "disk stars", "the <SUP>13</SUP>C pocket", "shape", "the same <SUP>13</SUP>C profile", "a flat <SUP>13</SUP>C profile", "gt", "C(α, n)<SUP>16</SUP>O reaction", "A" ]	8.199903	9.30058	-1
483926	[ "Kreckel, K.", "Armus, L.", "Groves, B.", "Lyubenova, M.", "Díaz-Santos, T.", "Schinnerer, E.", "Appleton, P.", "Croxall, K. V.", "Dale, D. A.", "Hunt, L. K.", "Beirão, P.", "Bolatto, A. D.", "Calzetti, D.", "Donovan Meyer, J.", "Draine, B. T.", "Hinz, J.", "Kennicutt, R. C.", "Meidt, S.", "Murphy, E. J.", "Smith, J. D. T.", "Tabatabaei, F. S.", "Walter, F." ]	2014ApJ...790...26K	[ "A Far-IR View of the Starburst-driven Superwind in NGC 2146" ]	16	[ "Max Planck Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany", "Spitzer Science Center, California Institute of Technology, MC 314-6, Pasadena, CA 91125, USA", "Max Planck Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany", "Max Planck Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany", "Spitzer Science Center, California Institute of Technology, MC 314-6, Pasadena, CA 91125, USA", "Max Planck Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany", "NASA Herschel Science Center, IPAC, California Institute of Technology, Pasadena, CA 91125, USA", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA", "Department of Physics and Astronomy, University of Wyoming, Laramie, WY 82071, USA", "INAF-Osservatorio Astrosico di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy", "Observatoire de Paris, 61 avenue de l'Observatoire, F-75014 Paris, France", "Department of Astronomy and Joint Space-Science Institute, University of Maryland, College Park, MD, USA; Max-Planck Institute for Astronomy, D-69117 Heidelberg, Germany; Visiting Humboldt Fellow.", "Department of Astronomy, University of Massachusetts, Amherst, MA 01003, USA", "National Radio Astronomy Observatory, Charlottesville, VA 22901, USA", "Princeton University Observatory, Peyton Hall, Princeton, NJ 08544-1001, USA", "Steward Observatory, University of Arizona, Tucson, AZ 85721, USA; MMT Observatory, Tucson, AZ 85721, USA", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Max Planck Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany", "Infrared Processing and Analysis Center, California Institute of Technology, MC 220-6, Pasadena, CA 91125, USA", "Department of Physics and Astronomy, University of Toledo, Toledo, OH 43606, USA", "Max Planck Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany", "Max Planck Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany" ]	[ "2014ApJ...794...26T", "2014PhyU...57.1199K", "2015A&A...574A..14C", "2015A&A...576A..33H", "2015ApJ...814...83L", "2015IAUS..309..322K", "2016ApJ...822...43J", "2017ApJ...844..155T", "2017MNRAS.465.2540P", "2018A&A...620A.140T", "2018Galax...6..138R", "2018MNRAS.474.3727L", "2019A&A...621A..51H", "2019A&A...631A..78S", "2020A&ARv..28....2V", "2023ApJ...953..113L" ]	[ "astronomy" ]	20	[ "dust", "extinction", "galaxies: individual: NGC 2146", "galaxies: starburst", "infrared: ISM", "ISM: kinematics and dynamics", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]	[ "1976A&A....53..397F", "1981PASP...93....5B", "1984ApJ...286..144W", "1987AJ.....93..264M", "1987AJ.....94..831A", "1988ApJ...325...74S", "1988ApJ...331L..81Y", "1988ngc..book.....T", "1989ApJ...347..727A", "1990AJ....100...60H", "1990ApJS...74..833H", "1991rc3..book.....D", "1992ApJ...390L..53K", "1995A&A...293..703M", "1995ApJ...445..666A", "1996AJ....111.1025M", "1997ApJS..112..315H", "1998ApJ...493..129S", "1999ApJ...514..772D", "1999ApJ...527..795K", "2000A&A...364..409G", "2000AJ....119.1608K", "2000ApJ...533..682C", "2001A&A...365..360T", "2001ApJ...554..981P", "2001ApJ...556..121K", "2001ApJ...561..766M", "2001ApJ...563..527G", "2002ApJ...580L..21W", "2003AJ....126.1607S", "2003ApJ...588...65S", "2003MNRAS.342..345C", "2003MNRAS.344.1000B", "2003MNRAS.346.1055K", "2004A&A...418L..27B", "2004AJ....127.2031K", "2004ApJ...604..534S", "2004ApJ...605..272S", "2004ApJ...613..898T", "2004ApJS..152..251V", "2004MNRAS.351.1151B", "2004PASP..116..138C", "2004SPIE.5487...50R", "2005ARA&A..43..769V", "2005ApJ...621..227M", "2005ApJ...621L..89B", "2005ApJ...630...82P", "2005ApJ...632..169L", "2005ApJS..160..115R", "2005PASJ...57..135I", "2005PASP..117..620R", "2006A&A...459..441G", "2006ApJ...642L.127E", "2006MNRAS.366.1151S", "2006PASP..118..129K", "2008A&A...479..687A", "2008A&A...492...31D", "2008ApJS..178...20A", "2008MNRAS.390..466E", "2009A&A...506.1541A", "2009AJ....137.4424T", "2009ApJ...692..187W", "2009ApJ...696..192W", "2009ApJ...707..250M", "2009ApJS..182..216K", "2009ApJS..184..230B", "2009PASJ...61..237T", "2009PASP..121..559A", "2010A&A...515A..35S", "2010A&A...517A..28M", "2010A&A...518L...1P", "2010A&A...518L...2P", "2010A&A...518L..41F", "2010A&A...518L..60B", "2010A&A...518L..66R", "2010A&A...522A...7A", "2010ApJ...711..818S", "2010ApJ...714L.162H", "2010ApJ...719.1503R", "2010MNRAS.401.2113E", "2010NewA...15..614G", "2011AJ....141...22S", "2011ApJ...730...28P", "2011ApJ...733L..16S", "2011ApJ...738...89S", "2011MNRAS.414..888E", "2011PASP..123.1218A", "2011PASP..123.1347K", "2012ApJ...749L..17Y", "2012ApJ...749L..18H", "2012MNRAS.425L..66M", "2013A&A...549A.118C", "2013A&A...557A..59B", "2013ApJ...763..120C", "2013ApJ...771...62K", "2013ApJ...775..127S", "2013ApJ...776...27V", "2013ApJ...777...66A", "2013ApJ...777...96C" ]	[ "10.1088/0004-637X/790/1/26", "10.48550/arXiv.1403.2381" ]	1403	1403.2381.txt	Galaxy outflows driven by star formation and active galactic nuclei (AGN) allow redistribution of energy and metals within the interstellar medium \citep{Veilleux2005} and have the potential to strongly affect galaxy evolution through quenching of star formation \citep{Kormendy2009}. They are a common feature in luminous infrared galaxies (LIRGs) \citep{Rupke2005}, which are observed to increasingly dominate the star formation rate density at higher redshifts up to z $\sim$2 \citep{LeFloch2005, PerezGonzalez2005}. NGC~2146 is one of the closest (17.2 Mpc, \citealt{Tully1988}) infrared luminous (L$_{\rm IR} = 1.2 \times$ 10$^{11}$ L$_\sun$, \citealt{Sanders2003}) galaxies. With a stellar mass of $2 \times 10^{10}$ M$_\sun$ \citep{Skibba2011} and morphological type SB(s)ab pec \citep{deVaucouleurs1991}, it displays a disturbed optical morphology due to a merger and a bright central bulge, extended irregular spiral arms, and deep dust lane features \citep{Greve2006}. H~\textsc{i} imaging reveals an extremely extended ($\sim$200 kpc) tail \citep{Fisher1976, Taramopoulos2001}. NGC~2146 has a well established superwind along the minor axis that has been detected in X-rays and optical emission lines from ionized atomic gas that show evidence for shock excitation \citep{Armus1995, Greve2000} as well as in a molecular gas outflow and superbubbles \citep{Tsai2009}. There is no evidence for AGN activity in the optical \citep{Ho1997} or in mid-IR IRS spectra \citep{Bernard-Salas2009, Petric2011}, though it may host a low-luminosity AGN based on its compact nuclear X-ray emission \citep{Inui2005}. As it exhibits a relatively high star formation rate of 7.9 M$_\sun$ year$^{-1}$ \citep{Kennicutt2011}, stellar winds and supernovae are thought to drive the superwind, similar to that observed in the starburst galaxy M82 \citep{McCarthy1987, Armus1989, Heckman1990}. However, as NGC~2146 is both larger and more massive compared to M82, its outflow exhibits a larger geometric scale and slower wind velocities \citep{Greve2000}. \begin{figure}[!ht] \centering \includegraphics[width=6.5in]{f1.pdf} \caption{R-band (left) and Spitzer/IRAC 3.6~$\mu$m (right) images of NGC~2146. Both are shown with a log scaling to show low surface brightness features. Overlaid are the field of view for the PPAK (white) and PACS spectroscopy (red) instruments, as well as an ellipse (blue) in order to provide a common reference for the remaining figures in this paper, with the center as determined from CO kinematic modeling \citep{Tsai2009} marked with an x. The ellipse corresponds to a diameter of 6 kpc at the assumed distance. \label{fig:fov}} \end{figure} Extensive analysis of the ionized and molecular gas kinematics have resulted in an understanding of the geometry of the system \citep{Greve2000, Tsai2009}. The disk is highly inclined at 63$^\circ$ \citep{dellaCeca1999}, with the near side of the disk in the south-west, making the south-west outflow behind and the north-east outflow in front of the disk. The outflow follows a conical morphology above and below the disk. The molecular gas outflow is seen through a breakout on the north-east side and two superbubbles, one to the south-east and another in projection along the major axis \citep{Tsai2009}. The combination of disk inclination angle and cone opening angle results in an unusual line-of-sight projection of the system, as diagramed in \cite{Greve2000} and \cite{Tsai2009}. The far side of the cone to the north and the near side of the cone to the south are essentially perpendicular to the line of sight, making their velocities consistent with the systemic velocity. The other walls of the cone fall, in projection, over the bright central disk component and may not be easily seen. Studies of nearby LIRGs and ULIRGs show they are generally interacting systems \citep{Armus1987, Sanders1988, Murphy1996}, and integral field spectroscopic surveys have revealed complex kinematics in multiple gas phases \citep{Arribas2008, Alonso-Herrero2009, Alonso-Herrero2010}. These kinematics develop along a merger sequence towards more elliptical or lenticular dynamics \citep{Bellocchi2013}. Many show evidence for shock ionization \citep{Armus1989, Monreal-Ibero2010}. However, optical studies are restricted by dust extinction, particularly along the galaxy disk and in the center where these winds are launched. Far-IR observations are able to penetrate these dusty regions, providing a direct view of the complete galactic wind. Such outflows have been observed within ULIRGs by molecular absorption against the nucleus \citep{Fischer2010, Sturm2011, Veilleux2013, Spoon2013}, but the possibility also exists to map outflow regions through the far-IR fine-structure emission lines. We present here results from Herschel PACS spectroscopy, which reveal conical outflows in the atomic and ionized gas that can be traced back to the central region of the galaxy disk. We link this to new optical integral field spectroscopy that traces the ionized shock diagnostics over a broad region both above and below the galaxy disk. We also examine the warm and cold dust distribution through multi-wavelength far-IR imaging. We present our observations in Section \ref{sec:data}, and our results in Section \ref{sec:results}. We discuss in Section \ref{sec:discussion} and conclude in Section \ref{sec:conclusion}. %>>>>>>>>>>>>>>>>>>>>>>>>>>>>>>>>>Section<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<<	\label{sec:conclusion} We present Herschel PACS far-IR fine structure line observations, optical IFU data and SOFIA 37 $\mu$m observations of the starburst driven superwind in NGC~2146, a nearby LIRG. We find high velocity dispersions in all far-IR lines, with deconvolved linewidths of $\sim$250~\kms~that extend along the minor axis through the disk center and open into conical regions above and below the galaxy disk. This is fully consistent with the previously proposed picture of the outflow geometry that results in an alignment of the far side of the cone to the north and the near side of the cone to the south perpendicular to the line of sight \citep{Greve2000, Tsai2009}, with the line emission from the opposing cone walls overwhelmed in projection by the bright central disk. Unlike in optical studies, which are limited by the high extinction at the galaxy center, we are able to observe the superwind launching region, which extends less than kpc along the major axis in the center of the galaxy. We present evidence for enhanced low-ionization line emission ([OI]/H$\alpha$, [SII]/H$\alpha$, [NII]/H$\alpha$) coincident with the regions of increased velocity dispersion above and below the plane seen in the far-IR emission line maps. The optical line ratios are consistent with shock excitation, though we cannot rule out a contribution from other low level ionization sources. The position of these shock and superwind indicators seen in the far-IR and optical have positions that match soft X-ray emission \citep{Inui2005}, with both the X-ray emission and CO bubbles \citep{Tsai2009} located within the full extent of the outflow cone we identify. We do not detect dust entrainment in the outflow, in either the warm or the cold dust components. In general, the outflow geometry is very similar to that observed in M82. Outflows of this nature would not be easily discerned in unresolved [CII] observations of high redshift systems but will be resolved by ALMA at z$\sim$1-3. We also observe stellar kinematics that are decoupled from the gas kinematics in all phases, consistent with predictions for post-merger systems, though we may also be suffering from some bias due to dust. As the galaxy is still in the midst of its post-merger starburst it has yet to expel or transform the bulk of its molecular gas. The starburst driven wind is potentially crucial for transformation of the galaxy's morphology into a red and dead elliptical. Far-IR line observations provide an unobscured view of the outflow and starburst, crucial for detailed study of the physical conditions in these regions. This system will provide an important local analog for future high redshift studies of IR-luminous galaxies driving galactic-scale superwinds.	14	3	1403.2381	NGC 2146, a nearby luminous infrared galaxy, presents evidence for outflows along the disk minor axis in all gas phases (ionized, neutral atomic, and molecular). We present an analysis of the multi-phase, starburst-driven superwind in the central 5 kpc as traced in spatially resolved spectral line observations, using far-IR Herschel PACS spectroscopy, to probe the effects on the atomic and ionized gas, and optical integral field spectroscopy to examine the ionized gas through diagnostic line ratios. We observe an increased ~250 km s<SUP>-1</SUP> velocity dispersion in the [O I] 63 μm, [O III] 88 μm, [N II] 122 μm, and [C II] 158 μm fine-structure lines that is spatially coincident with high excitation gas above and below the disk. We model this with a slow ~200 km s<SUP>-1</SUP> shock and trace the superwind to the edge of our field of view 2.5 kpc above the disk. We present new SOFIA 37 μm observations to explore the warm dust distribution, and detect no clear dust entrainment in the outflow. The stellar kinematics appear decoupled from the regular disk rotation seen in all gas phases, consistent with a recent merger event disrupting the system. We consider the role of the superwind in the evolution of NGC 2146 and speculate on the evolutionary future of the system. Our observations of NGC 2146 in the far-IR allow an unobscured view of the wind, crucial for tracing the superwind to the launching region at the disk center, and provide a local analog for future ALMA observations of outflows in high-redshift systems.	false	[ "high excitation gas", "diagnostic line ratios", "future ALMA observations", "optical integral field", "outflows", "high-redshift systems", "spatially resolved spectral line observations", "all gas phases", "the ionized gas", "the disk minor axis", "the regular disk rotation", "view", "new SOFIA 37 μm observations", "the disk center", "the system", "the atomic and ionized gas", "no clear dust entrainment", "a recent merger event", "the warm dust distribution", "the disk" ]	12.337517	8.556713	186
482822	[ "Burgess, J. Michael", "Preece, Robert D.", "Ryde, Felix", "Veres, Peter", "Mészáros, Peter", "Connaughton, Valerie", "Briggs, Michael", "Pe'er, Asaf", "Iyyani, Shabnam", "Goldstein, Adam", "Axelsson, Magnus", "Baring, Matthew G.", "Bhat, P. N.", "Byrne, David", "Fitzpatrick, Gerard", "Foley, Suzanne", "Kocevski, Daniel", "Omodei, Nicola", "Paciesas, William S.", "Pelassa, Veronique", "Kouveliotou, Chryssa", "Xiong, Shaolin", "Yu, Hoi-Fung", "Zhang, Binbin", "Zhu, Sylvia" ]	2014ApJ...784L..43B	[ "An Observed Correlation between Thermal and Non-thermal Emission in Gamma-Ray Bursts" ]	29	[ "Department of Space Science, University of Alabama in Huntsville, Huntsville, AL 35899, USA; Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA; The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova, SE-106 91 Stockholm, Sweden;", "Department of Space Science, University of Alabama in Huntsville, Huntsville, AL 35899, USA; Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA;", "Department of Physics, Royal Institute of Technology (KTH), AlbaNova, SE-106 91 Stockholm, Sweden; The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova, SE-106 91 Stockholm, Sweden;", "Department of Astronomy and Astrophysics, Pennsylvania State University, University Park, PA 16802, USA", "Department of Astronomy and Astrophysics, Pennsylvania State University, University Park, PA 16802, USA", "Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA", "Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA", "Physics Department, University College Cork, Cork, Ireland", "The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova, SE-106 91 Stockholm, Sweden; Department of Physics, Royal Institute of Technology (KTH), AlbaNova, SE-106 91 Stockholm, Sweden; Department of Physics, Stockholm University, AlbaNova, SE-106 91 Stockholm, Sweden", "Space Science Office, VP62, NASA/Marshall Space Flight Center, Huntsville, AL 35812, USA", "Department of Physics, Royal Institute of Technology (KTH), AlbaNova, SE-106 91 Stockholm, Sweden; The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova, SE-106 91 Stockholm, Sweden; Department of Astronomy, Stockholm University, SE-106 91 Stockholm, Sweden", "Department of Physics and Astronomy, Rice University, MS-108, P.O. Box 1892, Houston, TX 77251, USA", "Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA", "University College Dublin, Belfield, Dublin 4, Ireland", "University College Dublin, Belfield, Dublin 4, Ireland", "University College Dublin, Belfield, Dublin 4, Ireland; Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany", "W. W. Hansen Experimental Physics Laboratory, Kavli Institute for Particle Astrophysics and Cosmology, Department of Physics and SLAC National Accelerator Laboratory, Stanford University, Stanford, CA 94305, USA", "W. W. Hansen Experimental Physics Laboratory, Kavli Institute for Particle Astrophysics and Cosmology, Department of Physics and SLAC National Accelerator Laboratory, Stanford University, Stanford, CA 94305, USA", "Universities Space Research Association, Huntsville, AL 35805, USA", "Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA", "Department of Physics, Stockholm University, AlbaNova, SE-106 91 Stockholm, Sweden", "Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany", "Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA", "Department of Physics and Department of Astronomy, University of Maryland, College Park, MD 20742, USA" ]	[ "2014ApJ...795..155P", "2014ExA....38..331B", "2014NuPhS.256..241M", "2015A&A...573A..81Y", "2015A&C....12..100D", "2015ApJ...800L..23Z", "2015MNRAS.454L..31A", "2015PhR...561....1K", "2016A&A...588A.135Y", "2016ApJ...825...97U", "2016ApJS..225...17J", "2016MNRAS.460.1680I", "2016SGeo...37..705H", "2017NewA...55...22B", "2017SSRv..207...87B", "2018AN....339..352R", "2018ApJ...869..100U", "2018MNRAS.475..306R", "2018jwpw.book...87B", "2019ApJ...874...70C", "2019ApJ...886...20Y", "2019MNRAS.484.1912R", "2019MNRAS.485..474A", "2019arXiv190107505W", "2021ApJ...913...60P", "2021Galax...9...98L", "2021MNRAS.501.4974V", "2023arXiv230100576L", "2024ApJ...962...85Y" ]	[ "astronomy" ]	7	[ "gamma-ray burst: general", "radiation mechanisms: non-thermal", "radiation mechanisms: thermal", "Astrophysics - High Energy Astrophysical Phenomena" ]	[ "1979rpa..book.....R", "1981Natur.290..378M", "1982Natur.297..665F", "1985ApJ...288L..37M", "1986ApJ...308L..43P", "1986ApJ...308L..47G", "1993ApJ...413..281B", "1993ApJ...415..181M", "1996ApJ...466..768T", "1998ApJ...497L..17S", "2001ApJ...555..540L", "2002A&A...387..714D", "2002A&A...391.1141D", "2003ApJ...591..366K", "2006ApJS..166..298K", "2006RPPh...69.2259M", "2007ApJ...664L...1P", "2007ApJ...666.1012T", "2008GCN..8723....1W", "2009ApJ...697.1071A", "2009ApJ...702..791M", "2009ApJ...702.1211R", "2009GCN..8792....1V", "2010ApJ...709L.172R", "2011ApJ...726...90Z", "2011ApJ...727L..33G", "2011ApJ...730..141Z", "2011ApJ...732...34L", "2011ApJ...733L..40M", "2011ApJ...741...24B", "2012ApJ...755L...6F", "2012ApJ...757L..31A", "2012ApJS..199...19G", "2012MNRAS.419..573M", "2013A&A...551A.124H", "2013ApJ...764...94V", "2013ApJ...764..167S", "2013ApJ...770...32G", "2013MNRAS.428.2430L", "2013MNRAS.432.3237G", "2013MNRAS.433.2739I", "2014ApJ...784...17B", "2014Sci...343...42A", "2014Sci...343...51P" ]	[ "10.1088/2041-8205/784/2/L43", "10.48550/arXiv.1403.0374" ]	1403	1403.0374_arXiv.txt	Gamma-ray Bursts (GRBs) are believed to arise from the deaths of massive stars or the coalescence of two compact stellar objects such as neutron stars or black holes. The resulting explosion gives rise to an expanding fireball with a jet pointed at the observer but hidden from the observer until the density of radiation and particles in this highly relativistic outflow is low enough for radiation to escape, a region called the photosphere \citep[for a review see][]{meszaros:2006}. While the emission from this fireball is expected to be thermal \citep{Goodman:1986,Paczynski:1986}, observations over the past three decades suggest the prompt emission to be highly non-thermal \citep{Mazets:1981,Fenimore:1982,matz,Kaneko:2006,Goldstein:2012}, with only a few exceptions \citep{ryde:2010,ghirlanda:2013}. The conversion of the fireball energy into non-thermal $\gamma$-ray radiation involves the acceleration of electrons in the outflow and their subsequent cooling via an emission process such as synchrotron radiation \citep{sari:1998,tavani:1996}. Insight into these energy radiation emission processes in GRBs is obtained by comparing the observed $\gamma$-ray photon spectra directly to different radiation models. The $\fermi$ Gamma-ray Space Telescope offers a broad energy range for these comparisons. Recent observations \citep{Guiriec:2010,Zhang:2011,Axelsson:2012,guiriec:2013,Iyyani:2013,Preece:2013,Burgess:2013} show that at least two mechanisms can be present: a non-thermal component that is consistent with synchrotron emission from accelerated electrons in the jet and a typically smaller blackbody contribution from the photosphere. This photospheric emission is released when the fireball becomes optically thin so that an observer may see a mixture of thermal and non-thermal emission with different temporal characteristics that, when viewed together, can probe the development and structure of the fireball jet. This simple photospheric model has been used to quantitatively interpret several observed correlations such as the Amati correlation \citep[e.g.,][]{2007ApJ...666.1012T, 2011ApJ...732...34L, 2012ApJ...755L...6F} We are thus motivated to investigate correlations among spectral parameters derived by fitting the non-thermal component with a synchrotron photon model and the thermal component with a blackbody, an approach developed in previous investigations \citep{Burgess:2012,Burgess:2013}. The synchrotron model consists of an accelerated electron distribution, containing a relativistic Maxwellian and a high-energy power law tail that is convolved with the standard synchrotron kernel \citep{Burgess:2013,Burgess:2012,rybicki:1979}. We find that the characteristic energies ($E_{\rm p}$ for synchrotron and $kT$ for the blackbody) of the synchrotron and blackbody components are highly correlated across all the GRBs in our sample. We show that this correlation can be used to address the key question of how the energy of the outflow is distributed, i.e., whether the energy is in a magnetic field or is imparted as kinetic energy to baryons in the jet, and how this energy distribution evolves with time. %	The analysis of GRBs in the framework of this model can indicate whether the photosphere is in the acceleration or coasting phase, which in turn can be translated to the composition of the jet. We find that for exponents close to 2 the jet dynamics are dominated by the magnetic field while exponents close to 1 indicate baryonic jets. In our sample of six GRBs observed with $\fermi$, the exponents $\alpha$ of the relation between the characteristic energies of non-thermal and thermal components (Table \ref{tab:tab1}) span the range of possible values, showing that energy content of GRB jets ranges from being dominated by the magnetic field to being contained mostly in the kinetic energy of baryons in the jet. \edithere{A possible validation of this interpretation would be the future measurement of polarization in GRBs which will allow for the direct determination of the magnetization of GRB jets \citep[see for example][]{lundman:2013}.} \edithere{ We note that the lack of a correlation between the ratio of the thermal flux to the total flux with the inferred magnetic content of the jet is puzzling (see Table \ref{tab:tab1}). Naively, it is expected that a photosphere occurring deep in the acceleration phase of the outflow will have its thermal emission be much brighter than the non-thermal emission. A possible explanation for the weakness of the observed thermal component has been addressed by several authors \citep{zhang:2011,daigne:2013}. These works consider the effect of the magnetization parameter ($\sigma= \frac{B^2}{4\pi\Gamma\rho c^2}$) on the intensity of the thermal component where $\rho$ is the matter density of the outflow. For $\sigma\gg1$, most of the jet internal energy remains in the advected magnetic field, reducing the intensity of the observed thermal component from the photosphere. Another possibility for explaining the lack of correlation of the thermal flux ratios to the different jet modes is to consider that if the non-thermal flux is due to synchrotron following reconnection events above the photosphere, the amount of reconnection may not be simply given by the amount of magnetic energy and by the radius, but may depend also on the degree of tangledness of the field at that radius. For reconnection one needs field lines of opposite polarity near each other, and if the degree of randomness is stochastic (as it probably is), this could introduce a randomness in the amount of non-thermal electrons accelerated as well as the synchrotron flux produced. However, time-dependent simulations of magnetically dominated outflows in GRBs are not advanced enough to accurately test these assumptions and therefore the reduced intensity of the thermal component is still open to interpretation. }	14	3	1403.0374	Recent observations by the Fermi Gamma-ray Space Telescope have confirmed the existence of thermal and non-thermal components in the prompt photon spectra of some gamma-ray bursts (GRBs). Through an analysis of six bright Fermi GRBs, we have discovered a correlation between the observed photospheric and non-thermal γ-ray emission components of several GRBs using a physical model that has previously been shown to be a good fit to the Fermi data. From the spectral parameters of these fits we find that the characteristic energies, E <SUB>p</SUB> and kT, of these two components are correlated via the relation E <SUB>p</SUB>vpropT <SUP>α</SUP> which varies from GRB to GRB. We present an interpretation in which the value of the index α indicates whether the jet is dominated by kinetic or magnetic energy. To date, this jet composition parameter has been assumed in the modeling of GRB outflows rather than derived from the data.	false	[ "GRB", "several GRBs", "GRB outflows", "GRBs", "thermal and non-thermal components", "the observed photospheric and non-thermal γ-ray emission components", "Fermi", "SUP", "kT", "six bright Fermi GRBs", "the Fermi data", "E <SUB>p</SUB", "the prompt photon spectra", "the Fermi Gamma-ray Space Telescope", "kinetic or magnetic energy", "a good fit", "some gamma-ray bursts", "a physical model", "the characteristic energies", "the data" ]	16.172199	1.182188	5
800196	[ "Czerny, Michael", "Higaki, Tetsutaro", "Takahashi, Fuminobu" ]	2014PhLB..734..167C	[ "Multi-natural inflation in supergravity and BICEP2" ]	71	[ "Department of Physics, Tohoku University, Sendai 980-8578, Japan", "Theory Center, KEK, 1-1 Oho, Tsukuba, Ibaraki 305-0801, Japan", "Department of Physics, Tohoku University, Sendai 980-8578, Japan" ]	[ "2014JCAP...07..020C", "2014JCAP...08..001K", "2014JCAP...08..040D", "2014JCAP...10..025H", "2014JCAP...10..048G", "2014JCAP...10..053B", "2014JHEP...05..144C", "2014JHEP...07..074H", "2014PTEP.2014j3C01K", "2014PhLB..733..241C", "2014PhLB..734...68K", "2014PhLB..735..176C", "2014PhLB..736..526K", "2014PhLB..737..124K", "2014PhLB..738..301H", "2014PhLB..738..412G", "2014PhLB..738..457D", "2014PhRvD..90b3501L", "2014PhRvD..90b3506K", "2014PhRvD..90b3536C", "2014PhRvD..90d1303K", "2014PhRvD..90d3505B", "2014PhRvD..90d3519M", "2014PhRvD..90d7303C", "2014PhRvD..90h3503D", "2014SCPMA..57.1460C", "2015JCAP...09..006L", "2015JCAP...09..020R", "2015JHEP...02..127K", "2015JHEP...04..160A", "2015JHEP...10..062D", "2015JHEP...12..042B", "2015MPLA...3050106K", "2015PTEP.2015f3E02A", "2015PhLB..744..153H", "2015PhRvD..91h3516A", "2015PhRvD..92j3507P", "2015PhRvL.114q1301A", "2015arXiv150102354H", "2015arXiv150805335T", "2015arXiv151105885A", "2016JCAP...11..046K", "2016PhLB..757...32N", "2017AnPhy.376..460A", "2017JCAP...05..044D", "2017JCAP...07..011H", "2017PhLB..765..293D", "2017PhLB..773..179N", "2017arXiv170300453B", "2018JCAP...07..035L", "2018JHEP...02..104D", "2018PhRvD..98f1301B", "2019JCAP...06..047A", "2019JHEP...07..095T", "2019JHEP...08..029D", "2019JHEP...10..120T", "2021JCAP...04..007T", "2021JCAP...10..057T", "2021JHEP...01..152T", "2021PhRvD.104i5010K", "2022PhLB..83037143T", "2022PhRvD.105j3514J", "2022PhRvD.106e5014Y", "2023JCAP...12..039N", "2023JHEP...10..062M", "2023PDU....4201348P", "2023PhRvD.107f3016N", "2024JCAP...04..007F", "2024JCAP...05..060Y", "2024JHEP...05..327Y", "2024PhRvD.109c5024T" ]	[ "astronomy", "physics" ]	4	[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]	[ "1980PhLB...91...99S", "1981MNRAS.195..467S", "1981PhRvD..23..347G", "1982PhLB..108..389L", "1982PhRvL..48.1220A", "1983PhLB..129..177L", "1990PhRvL..65.3233F", "1993PhRvL..70.1912M", "1997PhRvL..78.1861L", "2000PhRvL..85.3572K", "2003PhLB..570..145F", "2003PhLB..570..151K", "2003PhRvD..68d6005K", "2004JHEP...04..003G", "2004JHEP...12..004K", "2005JCAP...01..005K", "2005JHEP...03..007B", "2005JHEP...08..007C", "2005NuPhB.718..113C", "2006JHEP...05..078C", "2006PhLB..636..126L", "2006PhLB..641..203K", "2006PhR...423...91G", "2007JHEP...01..103C", "2007JHEP...02..002K", "2007JHEP...02..028D", "2007PhR...445....1B", "2007PhRvD..75b5019A", "2007RvMP...79..733D", "2008JCAP...08..003D", "2008LNP...738..119K", "2008PhRvD..78j6003S", "2009JHEP...09..007B", "2009PhRvL.102l1301K", "2010JCAP...11..009N", "2010JCAP...11..039N", "2010PhLB..693..140T", "2010PhRvD..82d6003M", "2011JCAP...01..026K", "2011JCAP...02..010N", "2011JCAP...03..023K", "2011PhRvD..84d5021H", "2012JHEP...06..093H", "2012JHEP...07..195I", "2012JHEP...09..043H", "2012JHEP...11..125H", "2012NuPhB.858...63G", "2013JCAP...06..013T", "2013JCAP...08..031H", "2013JCAP...08..038N", "2013JHEP...07..005H", "2013NuPhB.874....1H", "2013PhLB..720..125H", "2013PhLB..724..165C", "2013PhLB..725..111N", "2013PhRvD..87d3520C", "2014JCAP...02..046J", "2014JCAP...05..046K", "2014JCAP...07..014C", "2014JCAP...08..012C", "2014JHEP...05..144C", "2014JHEP...09..140H", "2014JHEP...10..184A", "2014PhLB..730...24N", "2014PhLB..733..241C", "2014PhLB..734...96N", "2014PhLB..735..176C", "2014PhRvD..89j1302M", "2014PhRvD..90d3507G", "2014PhRvD..90h3003S", "2014PhRvL.112x1101B" ]	[ "10.1016/j.physletb.2014.05.041", "10.48550/arXiv.1403.5883" ]	1403	1403.5883_arXiv.txt	The BICEP2 experiment detected the primordial B-mode polarization~\cite{Ade:2014xna}, which could originate from the gravitational waves generated during inflation~\cite{Guth:1980zm,Linde:1981mu}. The tensor-to-scalar ratio $r$ and the Hubble parameter during inflation $H_{\rm inf}$ suggested by the BICEP2 are given by \bea \label{B-Hr} H_{\rm inf} &\simeq& 1.0 \times \GEV{14} \lrfp{r}{0.16}{\frac{1}{2}},\\ r &=& 0.20^{+0.07}_{-0.05} ~~(68\%{\rm CL}). \label{B-r} \eea The preferred range of $r$ is modified to $r = 0.16^{+0.06}_{-0.05}$ after subtracting the best available estimate for foreground dust. Taken at face value, the BICEP2 results strongly suggest large-field inflation such as chaotic inflation~\cite{Linde:1983gd}. For various large-field inflation models and their concrete realization in the standard model as well as supergravity and superstring theory, see e.g.~\cite{Freese:1990ni,Murayama:1992ua,Kawasaki:2000yn,Dimopoulos:2005ac,Kallosh:2007ig,Silverstein:2008sg,McAllister:2008hb,Kaloper:2008fb,Takahashi:2010ky, Nakayama:2010kt,Nakayama:2010sk,Harigaya:2012pg,Croon:2013ana,Nakayama:2013jka, Nakayama:2013nya,Cicoli:2014sva,Czerny:2014wza,Czerny:2014xja,Nakayama:2014-HCI}. It is worth noting, however, that the BICEP2 results are in tension with the {\it Planck} + WP + highL data. Explicitly, there is a tension on the relative size of scalar density perturbations on large and small scales. This tension suggests another extension of the Lambda CDM model such as a running spectral index, dark radiation, hot dark matter (HDM), non-zero neutrino mass~\cite{Ade:2014xna,Giusarma:2014zza}, anti-correlation between tensor and scalar modes~\cite{Contaldi:2014zua} or between tensor and isocurvature modes~\cite{Kawasaki:2014lqa}, and a sharp cut-off in the scalar modes~\cite{Miranda:2014wga}.\footnote{See e.g. Refs.~\cite{Feng:2003zua,Kawasaki:2003dd} for double inflation models to realize a sharp cut-off in the curvature perturbations.} If it is due to the running spectral index, it would provide us with invaluable information on the inflation sector.\footnote{ The implications of the other possibilities are also significant. The HDM candidates include axions or sterile neutrinos. The QCD axion can explain both hot and cold dark matter~\cite{Jeong:2013oza}. If both sterile neutrino HDM and dark radiation are present, there will be three coincidence problems of baryon-dark matter, neutrino-sterile neutrino, and photon-dark radiation densities. This may suggest the dark parallel world~\cite{Higaki:2013vuv}. } The BICEP2 data implies that the inflaton field excursion exceeds the Planck scale, $M_{pl} \simeq 2.4 \times \GEV{18}$~\cite{Lyth:1996im} . In order to have a sensible inflation model with super-Planckian field values, one has to control the inflaton potential over a broad field range. One way to accomplish this is to introduce a shift symmetry, under which the inflaton $\phi$ transforms as \bea \phi &\to& \phi +C, \eea where $\phi$ is the inflaton and $C$ is a real transformation parameter. As we shall see later, one of the plausible candidates for the inflaton with the above property is an axionic component of a moduli field. In fact, there appear many moduli fields through compactifications in string theory, and one of them could be the inflaton. The shift symmetry needs to be explicitly broken by some sources to generate the inflaton potential, and one plausible functional form is the sinusoidal function. In this case, the continuous shift symmetry is broken down to a discrete one. If a single sinusoidal function dominates the inflaton potential, it is the natural inflation~\cite{Freese:1990ni}. On the other hand, if there are many sources for the explicit breaking, it is the multi-natural inflation proposed in Ref.~\cite{Czerny:2014wza}, and its realization in supergravity was given in Ref.~\cite{Czerny:2014xja}. In multi-natural inflation, the inflaton potential consists of multiple sinusoidal functions with different height and periodicity: \bea \label{Vmulti} V_{\rm multi-natural} &=& \sum_i \Lambda_i^4 \cos\left(\frac{\phi}{f_i} + \theta_i \right) + {\rm const.}, \eea where the last constant term is such that the inflaton potential becomes zero at the potential minimum. Interestingly, only with two sinusoidal functions, a wide range of values of $(n_s, r)$ can be realized.\footnote{ Even small-field inflation is possible with two sinusoidal functions. See Ref.~\cite{Takahashi:2013tj} for curvatons with multiple sinusoidal functions. } We would like to emphasize the difference of the multi-natural inflation from $N$-flation~\cite{Dimopoulos:2005ac}.\footnote{ The realization of the $N$-flation in Large Volume Scenario~\cite{Balasubramanian:2005zx} was given in Ref.~\cite{Cicoli:2014sva}.} In the case of $N$-flation, there are many $(\gtrsim 100)$ axions, each of which has a sinusoidal potential. As a result, the high-scale inflation can be realized with sub-Planckian decay constants. In contrast, the multi-natural inflation is a single-field inflation, whose potential consists of multiple sinusoidal functions as in (\ref{Vmulti}), and the inflaton dynamics is much simpler. Recently, it was shown in Ref.~\cite{Czerny:2014wua} that multi-natural inflation has a built-in feature for generating a sizable running spectral index. This is because, if the inflaton potential has small modulations, a sizable running spectral index can be generated without spoiling the overall inflaton dynamics, as pointed out in Ref.~\cite{Kobayashi:2010pz}. The running spectral index, if confirmed by other experiments such as the 21cm observations~\cite{Shimabukuro:2014ava}, would tell us about rich structures of the inflaton potential. On the other hand, it is also possible that the tension between the BICEP2 and the Planck is due to some other sources. Even in this case, multi-natural inflation remains one of the viable large-field inflation models, which is one of the main targets of this paper. In this paper we revisit multi-natural inflation and its realization in supergravity in light of the recent BICEP2 results, and show that a wide range of values of $(n_s, r)$ can be realized. In particular, the spectral index close to or above unity with a large $r$ satisfying \EQ{B-r} is possible. Such large $n_s \sim 1$ will be favored if the tension between the BICEP2 and the Planck is solved by other sources such as dark radiation, HDM, or non-zero neutrino mass~\cite{Giusarma:2014zza}. We will also discuss the implications for the implementation of multi-natural inflation in string theory.	\label{Secsum} We have revisited multi-natural inflation~\cite{Czerny:2014wza} in supergravity for a UV completion. In this model the inflaton potential consists of two or more sinusoidal functions with different height and periodicity. A further UV completion based on a string-inspired framework has been also considered. The moduli stabilization of extra dimensions has to be done to give a sensible string theory in four dimensions. Multi-natural inflation may be realized as a remnant of such moduli stabilization based on various non-perturbative effects. From an observational point of view, multi-natural inflation is interesting because the existence of small modulations to the inflaton potential, which are necessary to generate a large running spectral index, is a built-in feature. The running spectral index is one way to relax the tension between BICEP2 and Planck. For this to work, we need a mild hierarchy between the two decay constants. On the other hand, the tension could be solved by other sources such as dark radiation\footnote{ The existence of dark radiation may be ubiquitous in the Large Volume Scenario~\cite{Cicoli:2012aq,Higaki:2013lra,Angus:2014bia}. }, in which case the scalar spectral index close to or even above unity is favored. We have shown that the multi-natural inflation with two sinusoidal functions can realize a wide range of values of $(n_s, r)$ that include such cases.	14	3	1403.5883	We revisit the recently proposed multi-natural inflation and its realization in supergravity in light of the BICEP2 results. Multi-natural inflation is a single-field inflation model where the inflaton potential consists of multiple sinusoidal functions, and it is known that a sizable running spectral index can be generated, which relaxes the tension between the BICEP2 and the Planck results. In this paper we show that multi-natural inflation can accommodate a wide range of values of (n<SUB>s</SUB>,r), including the spectral index close to or even above unity. This will be favored if the tension is resolved by other sources such as dark radiation, hot dark matter, or non-zero neutrino mass. We also discuss the implications for the implementation in string theory.	false	[ "Multi-natural inflation", "multi-natural inflation", "hot dark matter", "dark radiation", "multiple sinusoidal functions", "unity", "BICEP2", "a sizable running spectral index", "string theory", "non-zero neutrino mass", "other sources", "the BICEP2 results", "the spectral index", "s</SUB>,r", "values", "the recently proposed multi-natural inflation", "light", "non-zero", "supergravity", "the BICEP2 and the Planck results" ]	11.053406	-1.030105	-1
745802	[ "Li, Q.", "Liang, S. L.", "Li, A." ]	2014MNRAS.440L..56L	[ "Spectropolarimetric constraints on the nature of interstellar grains." ]	11	[ "Department of Physics, Xiangtan University, Xiangtan 411105, Hunan Province, China; Department of Physics and Astronomy, University of Missouri, Columbia, MO 65211, USA", "Department of Physics and Astronomy, University of Missouri, Columbia, MO 65211, USA", "Department of Physics and Astronomy, University of Missouri, Columbia, MO 65211, USA" ]	[ "2014MNRAS.445..301V", "2015ARA&A..53..541B", "2015ApJ...804...40G", "2016ApJ...825...22Y", "2016MNRAS.457.1593W", "2019ApJ...876...13H", "2020ApJ...892...96R", "2020ApJ...904...38Z", "2021ApJ...906...73H", "2021ApJ...909...94D", "2022dge..book.....W" ]	[ "astronomy" ]	4	[ "polarization", "dust", "extinction", "infrared: ISM", "Astrophysics - Astrophysics of Galaxies" ]	[ "1977ApJ...217..425M", "1984ApJ...285...89D", "1985ApJ...290..211L", "1987Natur.326..755K", "1988ApJ...333..848D", "1989ApJ...341..808M", "1989MNRAS.236..709D", "1990A&A...237..215D", "1990ASPC...12..193D", "1990QJRAS..31..567J", "1992A&A...259..614S", "1994ApJ...431..783K", "1994ApJ...433..797J", "1995ApJ...440..697S", "1995ApJ...455L.177G", "1996ApJ...467..334C", "1996ApJ...472..643M", "1997A&A...323..566L", "1998ApJ...497..824M", "1999ApJ...512..224A", "1999ApJ...524L..71M", "1999SSRv...90..149G", "2000MNRAS.312..327S", "2001A&A...367..355M", "2001ApJ...548..296W", "2001ApJ...554..778L", "2002AJ....124.2790I", "2002ApJ...577..789L", "2002ApJS..138...75P", "2002osp..conf...85W", "2003A&A...407L...5K", "2003ApJ...591.1220L", "2004ASPC..309..115W", "2004ASSL..319..535L", "2004ApJ...613..770M", "2004ApJS..152..211Z", "2005ASPC..343..293A", "2005ApJ...622..965L", "2006ApJ...651..268C", "2007A&A...463..635D", "2007ApJ...656..798M", "2007ApJ...657..810D", "2008ApJ...688L.103P", "2009ARA&A..47..481A", "2009ASPC..414..453D", "2009ApJ...696....1D", "2010A&A...513A..40K", "2010ARA&A..48...21H", "2010ApJ...718..867M", "2011AJ....141...22S", "2012A&A...539A.143N", "2012JQSRT.113.2334V", "2013A&A...558A..62J", "2013ApJ...770...78C", "2014A&A...561A..82S" ]	[ "10.1093/mnrasl/slu021", "10.48550/arXiv.1403.7138" ]	1403	1403.7138_arXiv.txt	} Although the exact nature of interstellar dust remains uncertain, it is now well recognized that interstellar grains consist of amorphous silicates and some form of carbonaceous materials. While the identification of amorphous silicate dust in the interstellar medium (ISM) is relatively secure through the broad featureless 9.7$\mum$ Si--O stretching and 18$\mum$ O--Si--O bending absorption features (see Henning 2010), our understanding of the carbon dust component, mainly through the 2175$\Angstrom$ extinction bump and the 3.4$\mum$ C--H absorption feature is not as clear. All contemporary dust models assume carbon dust as a key grain component. They differ mainly in terms of the exact chemical and physical forms the carbon dust component take: (1) The silicate-graphite model (Mathis et al.\ 1977; Draine \& Lee 1984; Siebenmorgen \& Kr\"ugel 1992; Weingartner \& Draine 2001; Draine \& Li 2007) assume that graphite is the major carbon sink and the silicate and graphite components are bare and physically separated; (2) The silicate core-carbonaceous mantle model (D\'{e}sert et al.\ 1990; Jones et al.\ 1990; Li \& Greenberg 1997; Jones et al.\ 2013) assumes that silicate grains are coated with a carbonaceous mantle made of either hydrogenated amorphous carbon (HAC) or organic refractory; (3) The composite model (Mathis \& Whiffen 1989; Mathis 1996; Zubko et al.\ 2004) assumes the dust to be low-density aggregates of small silicates and carbonaceous particles (amorphous carbon, HAC, and organic refractories). All dust models appear to be in general agreement with the observational constraints, including the interstellar extinction, scattering, polarization, IR emission and interstellar depletion. More recently, spectropolarimetry of the 3.4$\mum$ interstellar absorption feature has been used to distinguish between dust models (Adamson et al.\ 1999; Ishii et al.\ 2002; Chiar et al.\ 2006; Mason et al.\ 2007). The 3.4$\mum$ absorption feature, commonly attributed to the C--H stretching mode in saturated aliphatic hydrocarbon dust (see Pendleton \& Allamandola 2002), is ubiquitously seen in the diffuse ISM of the Milky Way and external galaxies (e.g., see Mason et al.\ 2004).\footnote{% Chiar et al.\ (2013) argued that the carrier of the 3.4$\mum$ feature is largely aromatic instead of aliphatic. But see Dartois et al.\ (2007) who argued that the carrier of the 3.4$\mum$ feature is highly aliphatic. } The silicate core-carbonaceous mantle model assumes that the 3.4$\mum$ absorption feature arises in the hydrocarbon mantles coating the amorphous silicate cores (Jones et al.\ 1990; Li \& Greenberg 1997). The hydrocarbon mantles consist of either ``organic refractory'' (Greenberg et al.\ 1995) or HAC (Jones et al.\ 1990). Both ``organic refractory'' and HAC provide a close match to the interstellar 3.4$\mum$ absorption feature (Greenberg et al.\ 1995; Mennella et al.\ 1999). % The interstellar organic refractory material is essentially HAC in character. The major difference between the organic refractory material with HAC lies in the way how they are made: the former is derived from the UV photo-processing of interstellar ice mixtures accreted on the pre-existing silicate cores (Greenberg et al.\ 1995); the latter results from direct accretion of gas-phase elemental carbon on the silicate cores in the diffuse ISM (Duley et al.\ 1989; Jones et al.\ 1990). The 9.7$\mum$ and 18$\mum$ silicate absorption features have been reported to be polarized along various sightlines probing both the diffuse ISM and dense molecular clouds (see Smith et al.\ 2000, Wright et al.\ 2002, Aitken 2005), suggesting that the silicate component is nonspherical and aligned. If the carrier of the 3.4$\mum$ feature resides in the carbonaceous mantles on the silicate cores, we would expect the 3.4$\mum$ absorption feature to be polarized as well (see Li \& Greenberg 2002). However, all spectropolarimetric observations show that the 3.4$\mum$ absorption feature is essentially unpolarized (Adamson et al.\ 1999, Ishii et al.\ 2002, Chiar et al.\ 2006, Mason et al.\ 2007). This, especially the nondetection of the 3.4$\mum$ feature polarization in the Galactic center Quintuplet combined with the fact that the 9.7$\mum$ silicate feature is polarized in the same sightline (Chiar et al.\ 2006), poses a severe challenge against the core-mantle model. They argue that the hydrocarbon dust do not reside on the same grains as the silicates, and likely form a separate population of small grains which are either spherical or not sufficiently aligned. This challenge might be alleviated if the mantle is much less elongated than the silicate core so that the 3.4$\mum$ feature would be polarized to a much smaller degree than the 9.7$\mum$ silicate feature. The aim of this work is to investigate this secanario. We consider an extreme case: the dust consists of a {\it spheroidal} silicate core coated with a layer of {\it spherical} aliphatic hydrocarbon mantle (\S2). To be complete, we also consider the composite model in \S3. We discuss these results in \S4. \vspace{-5mm}	} Since the spheroidal core-spherical mantle grains considered in \S2 represent an extreme case in which the 3.4$\mum$ hydrocarbon feature is least polarized (relative to the 9.7$\mum$ silicate feature), if in an astronomical object the 3.4$\mum$ absorption feature is observed to have an even lower degree of polarization, the core-mantle model will be severely challenged. Chiar et al.\ (2006) placed an upper limit on the 3.4$\mum$ feature polarization of the Galactic center Quintuplet object GCS\,3-II: $\PTOA \approx 0.13$. This upper limit is even lower than the lower limits predicted from the spheroidal core-spherical mantle model with $\vor/\vsil=2$ for which $\PTOA > 0.16$ (see Figure \ref{fig:P2A}d) over the entire allowable ranges of core-elongations ($1/3 < a_c/b_c < \sqrt{3}$; see \S2). This supports the idea put forward by Adamson et al.\ (1999) and Chiar et al.\ (2006) that, based on the nondetection of the 3.4$\mum$ feature polarization, the core-mantle model is invalid or the carrier of the 3.4$\mum$ feature does not reside in the carbonaceous mantle as previously thought. The hydrocarbon dust component responsible for the 3.4$\mum$ feature must be physically separated from the silicate component. This component must be either spherical or poorly aligned (or both) so that the resulting 3.4$\mum$ absorption feature is essentially unpolarized. With a {\it thicker} carbon mantle, one expects a {\it smaller} 3.4$\mum$ polarization-to-extinction ratio $\Pch/\Ach$. With $\vor/\vsil>3$, we obtain $\PTOA < 0.13$ which appears to satisfy the observational upper limit of Chiar et al.\ (2006). However, grains with such a thicker spherical mantle would produce little polarization in the optical wavelength range. This is inconsistent with the observations of the interstellar polarization: (1) light reaching us from reddened stars is often polarized in the optical; (2) the interstellar polarization curve --- the degree of polarization as a function of wavelength --- rises from the near-IR ($\lambda\sim 2\mum$), has a maximum somewhere in the optical ($\lambda_{\rm max}\approx 0.55\mum$) and then decreases toward the ultraviolet (UV; see Whittet 2003). To be considered successful, a grain model should have its bulk, submicrometer-sized dust component to be non-spherical and sufficiently aligned to reproduce the observed interstellar polarization curve (e.g., see Voshchinnikov 2012). The silicate-graphite model requires either silicate (e.g., Mathis et al.\ 1977, Kim \& Martin 1994) or both silicate and graphite (see Draine \& Fraisse 2009, Siebenmorgen et al.\ 2014) to account for the observed optical polarization. The core-mantle model requires the core-mantle dust to produce the interstellar optical polarization (see Li \& Greenberg 1997). The composite model requires the porous composite dust to account for the observed optical polarization (see Mathis \& Whiffen 1989, Mathis 1998). We note that, as shown in Figure \ref{fig:P2A}c, with $\vor/\vsil > 1$, the optical polarization-to-extinction ratio ($P_V/A_V$) predicted from the spheroidal core-spherical mantle model is too small to compare with the observational value of $P_V/A_V \le 0.064$ (Whittet 2003) which should be achieved for perfectly aligned grains. This indicates that, although with a thick spherical carbon mantle one may satisfy the observed upper limit of $\PTOA<0.13$, the starlight will essentially see the dust as {\it spherical} and will not be polarized in the optical. The composite model is not able to alleviate the 3.4$\mum$ polarization challenge. With $\PTOA \approx 1.02 - 1.08$ (compared with the observed upper limit of $\PTOA\approx 0.13$; Chiar et al.\ 2006), the composite model predicts a similar degree of polarization for the 9.7$\mum$ Si--O feature and the 3.4$\mum$ C--H feature. This indicates that the 3.4$\mum$ hydrocarbon feature should have a positive detection for the lines of sight along which the 9.7$\mum$ silicate feature is observed to be polarized. The nondetection of the 3.4$\mum$ feature polarization in the Galactic center Quintuplet combined with the detection of the 9.7$\mum$ silicate feature polarization in the same sightline (Chiar et al.\ 2006) poses a severe challenge against the composite model. The core-mantle model may remain valid if the mantle component does not contain the carrier of the 3.4$\mum$ absorption feature, i.e., the carrier of the 3.4$\mum$ absorption feature is not a major carbon sink in the ISM and is physically not associated with the bulk core-mantle dust. Jones et al.\ (2013) argued that the aliphatic hydrocarbon material is subject to UV photo-processing in the diffuse ISM and is expected to be maximally-aromatized in the order of a million years. Therefore, they suggested that the mantle material of the core-mantle dust is mainly aromatic and is not responsible for the 3.4$\mum$ absorption feature. According to Jones et al.\ (2013), the 3.4$\mum$ absorption feature is due to a separate population of small aliphatic hydrocarbon dust. If the core-mantle dust is not responsible for the 3.4$\mum$ absorption feature, we might encounter a carbon budget problem: if the total carbon abundance (relative to H) in the ISM is like the Sun $\cism=\csun\approx 224\ppm$ (Asplund et al.\ 2009)\footnote{% We note that the C abundance of the early B stars which are thought to be ideal indicators for the present-day interstellar abundances since they preserve their pristine abundances is close to the solar C abundance: $\cBstar\approx 214\pm20\ppm$ (Przybilla et al.\ 2008) and $\cBstar\approx 209\pm15\ppm$ (Nieva \& Przybilla 2012). } or proto-Sun $\cism=\csun\approx 288\ppm$ (Lodders 2003), with $\cgas\approx 140\ppm$ in the gas phase (Cardelli et al.\ 1996) and $\cpah\approx 60\ppm$ in PAHs (Li \& Draine 2001) subtracted, there is only $\cdust\approx 24\ppm$ or $\cdust\approx 57\ppm$ left for the 2175$\Angstrom$ extinction bump, the 3.4$\mum$ absorption feature, the ``extended red emission'' (ERE) which is most likely from some sort of small carbon-based dust (Witt \& Vijh 2004), and a population of bulk carbon dust.\footnote{% Sofia et al.\ (2011) derived $\cgas\approx 100 \ppm$ for several interstellar sightlines from the strong transition of C\,II] at 1334$\Angstrom$. They argued that the oscillator strength for the C\,II] transition at 2325$\Angstrom$ previously used by Cardelli et al.\ (1996) to obtain $\cgas\approx 140\ppm$ might have been underestimated. But even with $\cgas\approx 100 \ppm$, the amount of C available for the 2175$\Angstrom$ bump, the 3.4$\mum$ feature, the ERE, and the bulk carbon dust, is only $\cdust\approx 64\ppm$ or $\cdust\approx 97\ppm$ and does not seem to be sufficient. } The latter is required to account for part of the visual extinction since silicates alone are not able to provide enough extinction (see Footnote-14 in Li 2004). Furthermore, according to Draine (1990) and Jones et al.\ (1994), most of the dust mass in the ISM was condensed in the ISM, it is not very clear how it is possible for the re-condensation to keep the silicate and carbon grain populations apart in the ISM. Draine (2009) postulated a scenario for the ISM to grow two distinct grain types (i.e., silicate and carbon dust) out of a single gas mixture. He argued that when Mg, Si, Fe, and O atoms and ions arrive at the amorphous silicate surface, they are able to grow additional amorphous silicate; in contrast, the C atom physisorbed on the amorphous silicate surface might undergo photoexcitation to an excited state that is repulsive, ejecting it from the surface. Or perhaps the C would become hydrogenated or oxidized, with the resulting CH or CO undergoing photodesorption from the surface. Such processes could keep the amorphous silicate carbon-free in the diffuse ISM. Similar processes may occur on exposed carbonaceous surfaces: impinging C atoms could grow new carbonaceous material, whereas impinging Mg, Si, Fe atoms could be removed by some combination of reaction with impinging H or O, and photoexcitation by UV. \vspace{-5mm}	14	3	1403.7138	While it is well recognized that interstellar grains are made of amorphous silicates and some form of carbonaceous materials, it remains debated regarding what exact chemical and physical form the carbonaceous component takes. Contemporary grain models assume that the silicate and carbon components are either physically separated or they form a core-mantle structure, or they agglomerate to form porous composites. The core-mantle model posits that the mantle is made of some sort of aliphatic hydrocarbon materials and is responsible for the 3.4 μm absorption feature ubiquitously seen in the diffuse interstellar medium (ISM) of the Milky Way and external galaxies. This model is challenged by the non-detection of polarization in the 3.4 μm absorption feature as the 9.7 μm silicate feature is observed to be polarized. To alleviate this challenge, we calculate the degree of polarization of the 3.4 μm feature for spheroidal silicate dust coated by a layer of spherical aliphatic hydrocarbon. It is found that the 3.4 μm feature polarization still exceeds the observational upper limit, even though spherical aliphatic hydrocarbon mantles are expected to cause much less polarization than non-spherical (e.g. spheroidal) mantles. We have also shown that the composite grain model which consists of amorphous silicate, aliphatic hydrocarbon and vacuum also predicts the 3.4 μm feature polarization to well exceed what is observed. These results support the earlier arguments that the aliphatic hydrocarbon component is physically separated from the silicate component unless the 3.4 μm absorption feature is just a minor carbon sink in the ISM. L91	false	[ "spherical aliphatic hydrocarbon mantles", "spherical aliphatic hydrocarbon", "aliphatic hydrocarbon materials", "aliphatic hydrocarbon", "polarization", "carbonaceous materials", "the aliphatic hydrocarbon component", "porous composites", "silicate dust", "ISM", "amorphous silicate", "the 3.4 μm feature polarization", "amorphous silicates", "Contemporary grain models", "interstellar grains", "the 3.4 μm absorption feature", "the 9.7 μm silicate feature", "the carbonaceous component", "the silicate component", "the 3.4 μm feature" ]	10.919932	11.824379	-1

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