Datasets:
charlieoneill
/
astro_metadata

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12465084
[ "Eingorn, Maxim", "Kiefer, Claus", "Zhuk, Alexander" ]
2016JCAP...09..032E
[ "Scalar and vector perturbations in a universe with discrete and continuous matter sources" ]
24
[ "North Carolina Central University, CREST and NASA Research Centers, Fayetteville st. 1801, Durham, North Carolina 27707, U.S.A. ; Institute for Theoretical Physics, University of Cologne, Zülpicher Strasse 77, 50937 Köln, Germany;", "Institute for Theoretical Physics, University of Cologne, Zülpicher Strasse 77, 50937 Köln, Germany", "Astronomical Observatory, Odessa National University, Dvoryanskaya st. 2, Odessa 65082, Ukraine" ]
[ "2017ApJ...845..153B", "2017IJMPD..2643012E", "2017IJMPD..2650121E", "2017JCAP...12..022F", "2017PDU....17...63E", "2018EPJC...78..609A", "2018EPJC...78..637B", "2018Univ....4..109F", "2019EPJC...79..655E", "2019JCAP...10..065D", "2019PDU....2600329E", "2020PDU....2900565C", "2020PhLB..80935761B", "2021EPJC...81..246C", "2021EPJP..136..205E", "2021Univ....7..101E", "2021Univ....7..469E", "2022FrP....10.5757Z", "2022PhLB..82636911E", "2023EPJC...83..601B", "2023PhLB..83937797E", "2024JCAP...05..083E", "2024PhLB..85138564E", "2024arXiv240115214E" ]
[ "astronomy", "physics" ]
4
[ "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]
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[ "10.1088/1475-7516/2016/09/032", "10.48550/arXiv.1607.03394" ]
1607
1607.03394_arXiv.txt
In our paper we have studied a universe filled with dust-like matter in the form of discrete inhomogeneities (e.g., galaxies and their groups and clusters) which represents the CDM-component and additionally with two groups of other matter sources which can be responsible for dark energy. To cover a wide class of cases, we have considered a very general model where the first group of sources consists of perfect fluids with a linear EoS $p_I=\omega_I\varepsilon_I\, \ (\omega_I=\mathrm{const})$. The second group of matter sources consists of perfect fluids with an arbitrary nonlinear EoS: $p_J=f_J(\varepsilon_J)$. The background spacetime geometry is defined by the FLRW metric. We have developed the first-order scalar and vector cosmological perturbation theory. Our approach works at all cosmological scales and incorporates both linear and nonlinear effects with respect to energy density fluctuations. The only restriction is that we consider the weak gravitational field limit. We have demonstrated that the scalar perturbation $\Phi$ (i.e. the gravitational potential) as well as the vector perturbation $\mathbf{B}$ can be split into individual contributions from each matter source: $\Phi=\Phi_M+\sum_I\Phi_I+\sum_J\Phi_J$ and $\mathbf{B}=\mathbf{B}_M+\sum_I\mathbf{B}_I+\sum_J\mathbf{B}_J$. Each of these contributions satisfies its own equation (see Eqs. \rf{3.5}-\rf{3.7} and \rf{3.15}-\rf{3.17}). The velocity independent parts of $\Phi_M, \Phi_I$ and $\Phi_J$ are characterized by the finite time-dependent Yukawa interaction range $\lambda$, defined by the formula \rf{3.3} and being the same for each individual contribution. We have also obtained the exact form of $\Phi_M$ and $\mathbf{B}_M$ related to the discrete matter sources. We have performed a thorough check of the self-consistency of our approach. It is important to note that the equations obtained in our paper form the theoretical basis for subsequent numerical simulations for a very wide class of cosmological models. Since our approach is valid at arbitrary cosmological scales, we can use these equations for studying the mutual motion of galaxies and the Hubble flow formation at relatively small scales (e.g., up to 20-30 Mpc), as well as for the investigation of structure formation at very large cosmological distances 1000-3000 Mpc corresponding to the largest known cosmic structures \cite{supstr1,supstr2,supstr3}. The formation of such enormously large structures is a challenge of modern cosmology because they considerably exceed the previously reported cell of homogeneity dimension $\approx 370$ Mpc \cite{Yadav}.
16
7
1607.03394
We study a universe filled with dust-like matter in the form of discrete inhomogeneities (e.g., galaxies and their groups and clusters) and two sets of perfect fluids with linear and nonlinear equations of state, respectively. The background spacetime geometry is defined by the FLRW metric. In the weak gravitational field limit, we develop the first-order scalar and vector cosmological perturbation theory. Our approach works at all cosmological scales (i.e. sub-horizon and super-horizon ones) and incorporates linear and nonlinear effects with respect to energy density fluctuations. We demonstrate that the scalar perturbation (i.e. the gravitational potential) as well as the vector perturbation can be split into individual contributions from each matter source. Each of these contributions satisfies its own equation. The velocity-independent parts of the individual gravitational potentials are characterized by a finite time-dependent Yukawa interaction range being the same for each individual contribution. We also obtain the exact form of the gravitational potential and vector perturbation related to the discrete matter sources. The self-consistency of our approach is thoroughly checked. The derived equations can form the theoretical basis for numerical simulations for a wide class of cosmological models.
false
[ "energy density fluctuations", "individual contributions", "cosmological models", "super-horizon ones", "numerical simulations", "discrete inhomogeneities", "linear", "perfect fluids", "linear and nonlinear equations", "state", "linear and nonlinear effects", "the discrete matter sources", "respect", "the individual gravitational potentials", "-", "each matter source", "each individual contribution", "the first-order scalar and vector cosmological perturbation theory", "the gravitational potential and vector perturbation", "Yukawa" ]
10.654408
0.552533
89
2774928
[ "Accurso, G.", "Saintonge, A.", "Bisbas, T. G.", "Viti, S." ]
2017MNRAS.464.3315A
[ "Radiative transfer meets Bayesian statistics: where does a galaxy's [C II] emission come from?" ]
29
[ "Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK", "Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK", "Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany; Department of Astronomy, University of Florida, Gainesville, FL 32611, USA", "Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK" ]
[ "2017ApJ...840...51L", "2017ApJ...846...32D", "2017ApJ...846..105O", "2017MNRAS.470.4750A", "2018ApJ...869...73L", "2018MNRAS.481.1976Z", "2018MNRAS.481.4277F", "2018NewAR..82....1H", "2019A&A...626A..23C", "2020A&A...643A...5D", "2020A&A...643A...6C", "2020A&A...643A.141M", "2020ApJ...899...23H", "2020ApJ...903...30B", "2020RSOS....700556H", "2021A&A...651A..59L", "2021A&A...653L..10Z", "2021MNRAS.500.3802H", "2021MNRAS.502.2701B", "2022A&A...658A.151G", "2022A&A...666A.112L", "2022ApJ...929...92V", "2022ApJ...934..115B", "2023ApJ...950..119T", "2023MNRAS.518.3074M", "2023MNRAS.519..729B", "2024A&A...685A..80Z", "2024ApJ...961...42T", "2024MNRAS.527.8886B" ]
[ "astronomy" ]
7
[ "astrochemistry - ISM: molecules", "photodissociation region (PDR)", "ISM: structure", "infrared: galaxies", "infrared: ISM", "Astrophysics - Astrophysics of Galaxies" ]
[ "1968IAUS...34..205F", "1977A&A....55..137D", "1978ApJS...36..595D", "1979ApJS...41..555H", "1980A&A....91...68D", "1984ApJ...285...89D", "1988ApJ...334..771V", "1991ApJ...374..580B", "1991ApJ...377..192H", "1994A&AS..105...29F", "1994ApJ...423..223C", "1994ApJ...436..720H", "1996A&A...311..690L", "1997A&AS..121...15H", "1999A&A...344..282L", "1999ApJ...527..795K", "1999ApJS..123....3L", "2000A&A...353..276K", "2000ApJ...545L.121P", "2001MNRAS.322..231K", "2003MNRAS.340.1136E", "2003MNRAS.346.1055K", "2003PASP..115..763C", "2004ApJ...604..222C", "2004ApJ...613..898T", "2004MNRAS.348.1337W", "2004MNRAS.351.1151B", "2005ApJ...621..328H", "2005ApJ...621..695V", "2005ApJ...622..759G", "2005ApJS..161...65A", "2005MNRAS.362.1038E", "2006A&A...451..917R", "2006ApJ...652L.125O", "2006MNRAS.371.1865B", "2006PNAS..10312269D", "2007A&A...468...33E", "2007ApJ...660L..43N", "2007ApJ...665.1489K", "2007ApJ...670..156D", "2008MNRAS.385.2011P", "2009A&A...507.1327H", "2010A&A...521L..17L", "2010A&A...521L..18V", "2010ApJ...713..871W", "2010ApJ...720..226P", "2010ApJ...721..193P", "2010ApJ...724..957S", "2010ApJS..189..309L", "2010CAMCS...5...65G", "2010HiA....15..408V", "2010PASP..122..261B", "2010PhDT........91V", "2011A&A...529A.149G", "2011A&A...532A.152M", "2011ApJ...731..120M", "2011ApJ...737...12L", "2011ApJS..197...16W", "2011MNRAS.415...11D", "2011MNRAS.416.2712D", "2012A&A...540A..52C", "2012A&A...548A..20C", "2012ApJ...754...25R", "2012ApJS..203...13G", "2012MNRAS.420..141E", "2012MNRAS.421...98D", "2012MNRAS.427.2100B", "2013A&A...550A..36M", "2013A&A...550A.114B", "2013A&A...553A.114K", "2013A&A...554A.103P", "2013ARA&A..51..393C", "2013ApJ...779...46H", "2013MNRAS.428.1606F", "2013MNRAS.430..288S", "2013MNRAS.431.2493K", "2013RMxAA..49..137F", "2013ascl.soft03004V", "2014A&A...561A.122L", "2014AAS...22333604L", "2014ApJ...784....3C", "2014ApJ...792...34O", "2014ApJ...794...45M", "2014ApJ...796...84R", "2014ApJS..214...15S", "2014MNRAS.440..942C", "2014MNRAS.440L..81O", "2014MNRAS.443..111B", "2014MNRAS.444.3894H", "2015A&A...575A..17H", "2015A&A...578A..70K", "2015A&A...579A.102B", "2015MNRAS.452...54M", "2015MNRAS.454.2828B", "2016A&A...586A..37M", "2016ApJ...816...23S", "2016ApJ...816L..23F", "2016MNRAS.457.3306O", "2016MNRAS.457.3334G" ]
[ "10.1093/mnras/stw2580", "10.48550/arXiv.1607.03488" ]
1607
1607.03488_arXiv.txt
This is a simple template for authors to write new MNRAS papers. See \texttt{mnras\_sample.tex} for a more complex example, and \texttt{mnras\_guide.tex} for a full user guide. All papers should start with an Introduction section, which sets the work in context, cites relevant earlier studies in the field by \citet{Others2013}, and describes the problem the authors aim to solve \citep[e.g.][]{Author2012}.
The last numbered section should briefly summarise what has been done, and describe the final conclusions which the authors draw from their work.
16
7
1607.03488
The [C II] 158 μm emission line can arise in all phases of the interstellar medium (ISM), therefore being able to disentangle the different contributions is an important yet unresolved problem when undertaking galaxy-wide, integrated [C II] observations. We present a new multiphase 3D radiative transfer interface that couples STARBURST99, a stellar spectrophotometric code, with the photoionization and astrochemistry codes MOCASSIN and 3D-PDR. We model entire star-forming regions, including the ionized, atomic, and molecular phases of the ISM, and apply a Bayesian inference methodology to parametrize how the fraction of the [C II] emission originating from molecular regions, f_{[C II],mol}, varies as a function of typical integrated properties of galaxies in the local Universe. The main parameters responsible for the variations of f_{[C II],mol} are specific star formation rate (SSFR), gas phase metallicity, H II region electron number density (n<SUB>e</SUB>), and dust mass fraction. For example, f_{[C II],mol} can increase from 60 to 80 per cent when either n<SUB>e</SUB> increases from 10<SUP>1.5</SUP> to 10<SUP>2.5</SUP> cm<SUP>-3</SUP>, or SSFR decreases from 10<SUP>-9.6</SUP> to 10<SUP>-10.6</SUP> yr<SUP>-1</SUP>. Our model predicts for the Milky Way that f_{[C II],mol} = 75.8 ± 5.9 per cent, in agreement with the measured value of 75 per cent. When applying the new prescription to a complete sample of galaxies from the Herschel Reference Survey, we find that anywhere from 60 to 80 per cent of the total integrated [C II] emission arises from molecular regions.
false
[ "H II region electron number density", "molecular regions", "II", "dust mass fraction", "cent", "gas phase metallicity", "galaxies", "f_{[C II],mol", "typical integrated properties", "ISM", "f_{[C", "Universe", "specific star formation rate", "SSFR", "the total integrated [C II] emission", "158 μm emission line", "entire star-forming regions", "SUP>-9.6</SUP", "cm", ", integrated [C II] observations" ]
12.736025
8.552049
185
4924940
[ "Aravena, M.", "Decarli, R.", "Walter, F.", "Bouwens, R.", "Oesch, P. A.", "Carilli, C. L.", "Bauer, F. E.", "Da Cunha, E.", "Daddi, E.", "Gónzalez-López, J.", "Ivison, R. J.", "Riechers, D. A.", "Smail, I.", "Swinbank, A. M.", "Weiss, A.", "Anguita, T.", "Bacon, R.", "Bell, E.", "Bertoldi, F.", "Cortes, P.", "Cox, P.", "Hodge, J.", "Ibar, E.", "Inami, H.", "Infante, L.", "Karim, A.", "Magnelli, B.", "Ota, K.", "Popping, G.", "van der Werf, P.", "Wagg, J.", "Fudamoto, Y." ]
2016ApJ...833...71A
[ "The ALMA Spectroscopic Survey in the Hubble Ultra Deep Field: Search for [CII] Line and Dust Emission in 6" ]
99
[ "Núcleo de Astronomía, Facultad de Ingeniería, Universidad Diego Portales, Av. Ejército 441, Santiago, Chile", "Max-Planck Institut für Astronomie, Königstuhl 17, D-69117, Heidelberg, Germany", "Max-Planck Institut für Astronomie, Königstuhl 17, D-69117, Heidelberg, Germany; Astronomy Department, California Institute of Technology, MC105-24, Pasadena, CA 91125, USA ; NRAO, Pete V. Domenici Array Science Center, P.O. Box O, Socorro, NM 87801, USA", "Leiden Observatory, Leiden University, P.O. Box 9513, NL2300 RA Leiden, The Netherlands ; UCO/Lick Observatory, University of California, Santa Cruz, CA 95064, USA", "Astronomy Department, Yale University, New Haven, CT 06511, USA", "Leiden Observatory, Leiden University, P.O. Box 9513, NL2300 RA Leiden, The Netherlands ; Astrophysics Group, Cavendish Laboratory, J. J. Thomson Avenue, Cambridge CB3 0HE, UK", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile Av. Vicuña Mackenna 4860, 782-0436 Macul, Santiago, Chile ; Millennium Institute of Astrophysics, Chile ; Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO 80301, USA", "Research School of Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611, Australia ; Centre for Astrophysics and Supercomputing, Swinburne University of Technology, Hawthorn, Victoria 3122, Australia", "Laboratoire AIM, CEA/DSM-CNRS-Universite Paris Diderot, Irfu/Service d'Astrophysique, CEA Saclay, Orme des Merisiers, F-91191 Gif-sur-Yvette cedex, France", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile Av. Vicuña Mackenna 4860, 782-0436 Macul, Santiago, Chile", "European Southern Observatory, Karl-Schwarzschild Strasse 2, D-85748 Garching bei München, Germany ; Institute for Astronomy, University of Edinburgh, Blackford Hill, Edinburgh EH9 3HJ, UK", "Cornell University, 220 Space Sciences Building, Ithaca, NY 14853, USA", "Institute for Computational Cosmology, Durham University, South Road, Durham DH1 3LE, UK", "Institute for Computational Cosmology, Durham University, South Road, Durham DH1 3LE, UK", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany", "Millennium Institute of Astrophysics, Chile ; Departamento de Ciencias Físicas, Universidad Andres Bello, Fernandez Concha 700, Las Condes, Santiago, Chile", "Université Lyon 1, 9 Avenue Charles André, F-69561 Saint Genis Laval, France", "Department of Astronomy, University of Michigan, 500 Church Street, Ann Arbor, MI 48109, USA", "Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany", "NRAO, Pete V. Domenici Array Science Center, P.O. Box O, Socorro, NM 87801, USA ; Joint ALMA Observatory—ESO, Av. Alonso de Córdova, 3104, Santiago, Chile", "Joint ALMA Observatory—ESO, Av. Alonso de Córdova, 3104, Santiago, Chile", "Leiden Observatory, Leiden University, P.O. Box 9513, NL2300 RA Leiden, The Netherlands", "Instituto de Física y Astronomía, Universidad de Valparaiso, Avda. Gran Bretaña 1111, Valparaiso, Chile", "Université Lyon 1, 9 Avenue Charles André, F-69561 Saint Genis Laval, France", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile Av. Vicuña Mackenna 4860, 782-0436 Macul, Santiago, Chile", "Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany", "Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany", "Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK ; Cavendish Laboratory, University of Cambridge, 19 J.J. Thomson Avenue, Cambridge CB3 0HE, UK", "European Southern Observatory, Karl-Schwarzschild Strasse 2, D-85748 Garching bei München, Germany", "Leiden Observatory, Leiden University, P.O. Box 9513, NL2300 RA Leiden, The Netherlands", "SKA Organization, Lower Withington Macclesfield, Cheshire SK11 9DL, UK", "European Southern Observatory, Karl-Schwarzschild Strasse 2, D-85748 Garching bei München, Germany; Universität-Sternwarte München, Scheinerstr. 1, D-81679 München, Germany" ]
[ "2016ApJ...833...67W", "2016ApJ...833...68A", "2016ApJ...833...70D", "2016ApJ...833...72B", "2016ApJ...833...73C", "2016ApJ...833..153S", "2017A&A...605A..42C", "2017A&A...608A...1B", "2017A&A...608A.138G", "2017ApJ...834...36H", "2017ApJ...836....8T", "2017ApJ...845..108Y", "2017ApJ...847...21F", "2017ApJ...848...30W", "2017ApJ...848...49C", "2017ApJ...848...78F", "2017FrASS...4...49T", "2017MNRAS.465.3475W", "2017MNRAS.469.4863B", "2017MNRAS.472.4587H", "2017Natur.545..457D", "2017PASJ...69...45H", "2017PhyU...60..961S", "2017arXiv170909066K", "2018A&A...609A.130L", "2018A&A...617A..82B", "2018ASPC..517..557C", "2018ApJ...856..107S", "2018ApJ...861..100C", "2018ApJ...862...77C", "2018ApJ...862...78C", "2018ApJ...867..153C", "2018IAUS..333..209Y", "2018MNRAS.473.1909G", "2018MNRAS.474.1718N", "2018MNRAS.478.1170C", "2018Natur.553..178S", "2018Natur.556..469M", "2018PASA...35....2V", "2018arXiv180101508T", "2018arXiv181007536C", "2018arXiv181204283R", "2019A&A...621A.133B", "2019A&A...627A.131C", "2019A&A...632A.106L", "2019ApJ...882..136A", "2019ApJ...882..138D", "2019ApJ...882..139G", "2019ApJ...886...60S", "2019ApJ...887..107F", "2019IAUS..344..259L", "2019MNRAS.487.1844M", "2019MNRAS.487.1946Q", "2019MNRAS.488.3014P", "2019MNRAS.490.1928Y", "2019arXiv190308177L", "2019astro2020T.402K", "2020A&A...637A..32B", "2020ApJ...889...98M", "2020ApJ...895...74N", "2020ApJ...901...79A", "2020ApJ...902..110D", "2020IAUS..352..177A", "2020IAUS..352..216S", "2020MNRAS.494.1071G", "2020MNRAS.496..875R", "2020Msngr.179...17A", "2020RSOS....700556H", "2021ApJ...911...99F", "2021ApJ...911..132Y", "2021ApJ...912...62G", "2021ApJ...912...67U", "2021ApJ...915...33S", "2021ApJ...923..215C", "2021IAUS..356..261T", "2021MNRAS.500.3083C", "2021NatAs...5.1110W", "2022A&A...667A.156B", "2022ApJ...928...68S", "2022MNRAS.510.5603K", "2023A&A...673A.153P", "2023A&A...674A.161T", "2023A&A...679A.131R", "2023ApJ...943..151G", "2023ApJ...945..111B", "2023ApJ...958...16W", "2023MNRAS.521.6124P", "2023MNRAS.523.3503P", "2023Natur.617..261K", "2023PhyU...66.1071S", "2023arXiv230301658F", "2023arXiv230312676S", "2023arXiv230907834F", "2024A&A...682A..24A", "2024A&A...685A.121L", "2024ApJ...965..108U", "2024ApJ...968..113E", "2024arXiv240205849U", "2024arXiv240607512C" ]
[ "astronomy" ]
9
[ "galaxies: evolution", "galaxies: high-redshift", "galaxies: ISM", "galaxies: star formation", "instrumentation: interferometers", "submillimeter: galaxies", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.3847/1538-4357/833/1/71", "10.48550/arXiv.1607.06772" ]
1607
1607.06772.txt
A key to understanding galaxy formation is to investigate the physical mechanisms that lead to the formation of the first galaxies and thereby to understand their role in the reionization of the Universe \citep{robertson10}. One of the main challenges in studying galaxies within the first Gigayear of the Universe (i.e. $z>6$) is that observations in the optical and near-infrared (NIR) regimes can only probe the high-resonance Ly-$\alpha$ line (rest wavelength: 1216\AA) and the faint underlying UV continuum. Both measurements are strongly affected by dust obscuration, and the Ly-$\alpha$ line is known to be hard to interpret due to the difficulties of radiative transfer modelling. Despite significant observational efforts, detection of Ly-$\alpha$ emission in non--quasar environments at $z>6$ has been very scarce. One interpretation of this finding is that the increasingly neutral intergalactic medium (IGM) that is surrounding galaxies during the epoch of reionization absorbs the Ly-$\alpha$ emission line through its damping wings. This in turn severely limits its usefulness as a spectroscopic redshift indicator \citep[e.g.,][]{schenker12,pentericci14} Since the strength of the Ly-$\alpha$ line has been found to decline very rapidly beyond $z > 6$, optical spectroscopic confirmation has proven to be challenging with current facilities \citep[e.g.][]{treu13}. Moreover, at $z>6.5$, the Ly$\alpha$ line shifts into a range of the electromagnetic spectrum highly contaminated by sky lines, making the identification of $z\sim7$ sources even more challenging. Beyond that, the line enters the near--IR bands, that are limited in sensitivity through ground--based observations; this situation will not change until the launch of the {\it James Webb} Space Telescope ({\it JWST}) . As a consequence, only a handful of sources with spectroscopically confirmed redshift at $z > 6.5$ are known to date \citep[][]{oesch15}. The far-infrared (FIR) fine-structure [CII]158$\mu$m emission line has been proposed as an alternative to Ly-$\alpha$ to study the first galaxies at $z>6$ \citep[e.g.][]{walter12}. The [CII] line is the dominant coolant of the interstellar medium (ISM) in star forming galaxies, making up 0.1-1\% of the integrated FIR luminosity of galaxies \citep[e.g. early work by ][]{stacey91}. This line appears to be ubiquitous in star forming galaxies, and at $z>6$ is redshifted into the accessible 1-mm band observable from the ground. Several studies have recently shown the power of this line to study objects in this redshift range \citep{carilli13}, but the majority of sources detected at z$>$6 are quasar host galaxies \citep[e.g.][]{maiolino05,walter09,venemans12,wang13,venemans16}. The current highest-redshift detections of [CII] emission in non--AGN--dominated galaxies, using ALMA, are at $z = 5.7-6.3$ \citep{riechers13,capak15,willott15,gullberg15,maiolino16,knudsen16}. The detection of such systems was not possible in the pre--ALMA era given the collecting area of the previous generation of millimeter interferometers. The brightness of the [CII] line in principle makes it a unique tool to investigate the properties of galaxies well into the reionization epoch. Since this line is bright in typical star forming galaxies, it can be readily used as a direct way to identify and spectroscopically secure galaxies in blind millimeter spectroscopic surveys of the sky, in particular in galaxies that cannot be followed up spectroscopically in the optical/NIR (at least not before the launch if {\it JWST}). Compared to the more conventional tracer of the ISM, CO emission, the [CII] line is intrinsically much brighter. Given the rather high excitation temperature of $\sim$92\,K it also possible that the [CII] line is much less susceptible to the effects of the cosmic microwave background (CMB) than the CO emission \citep{dacunha13}, even though recent studies have suggested that in the cold neutral ISM, the spin temperature of the [CII] transition is almost coupled to the CMB temperature, and thus the [CII] luminosity could be as low as 20\% of the one that one could have inferred without taking into account of the CMB \citep{vallini15}. Using the Atacama Large Millimeter/submillimeter Array (ALMA), we have obtained the first full 1\,mm spectroscopic survey in a region of the cosmological deep field for which the deepest multi--wavelength data exist: the {\it Hubble} Ultra Deep Field \citep[UDF;][]{beckwith06}. This survey enables, for the first time, a blind search for [CII] emission in a redshift range $6<z<8$. The UDF is an ideal field to perform such a study, as this field contains a high density of dropout galaxies in this redshift range \citep[e.g.,][]{bouwens14,bouwens15,schenker12,mclure11, mclure13}. In this paper, we present the result of our search for [CII] line emission. This pathfinder study demonstrates that ALMA's tuning range and sensitivity is well--matched to detecting the normal galaxy population at $z>6$. The layout of this paper is as follows. In \S \ref{sect:obs}, we briefly describe our ALMA observations of the UDF and the multi--wavelength data of this field used in this study. In \S \ref{sect:results}, we introduce our methodology and algorithm to search for [CII] line emitters, present our list of candidate [CII] sources, and the level of contamination and completeness of our catalog. In this Section, we analyse the possibility that our line detections correspond to other redshift solutions based on the existence of other molecular line detections in our spectral scans. Here, we also present a blind [CII] line candidate based on possible detection of [CII] and CO(6-5) line emission. In \S \ref{sect:discussion}, we investigate the location of our [CII] line candidates in the SFR-[CII] luminosity plane and place first constraints on the [CII] luminosity function at $6<z<8$. Finally, in \S \ref{sect:conclusion}, we list the main conclusions of this study. We adopt a standard $\Lambda$CDM cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda}=0.7$ and $\Omega_{\rm M}=0.3$.
\label{sect:conclusion} We have used our molecular ALMA survey of the {\it Hubble} UDF to search for [CII] line candidates in the redshift range $6.0<z<8.0$. We do this by blindly searching for significant line peaks in the 1-mm data cube, and by specifically looking around 58 known dropout galaxies with photometric redshifts that are consistent with the redshift range covered by [CII] ($5.5<z<8.5$). Our survey field within the UDF was chosen such that the number of known drop--out sources was maximised. The spectroscopic survey reaches an approximately uniform depth of L$_{\rm [CII]}\sim$(1.6-2.5)$\times$10$^{8}$\,L$_{\odot}$ at an angular resolution of $\sim$\,1$"$, well matched to the expected size of the galaxies at that redshift. We discuss our statistical tools to search for line emission in the vicinity of the drop--out galaxies. These include assessing the {\it fidelity} (fraction of positive vs.\ negative emission line candidates) as well as the {\it completeness} (our ability to recover artificial objects) of the line search. We find that there are more positive line candidates than (physically implausible) negative ones, assuring us that we are recovering actual signal in our spectroscopic survey at the positions of interest. We end up with 14 [CII] line candidates above a S/N cutoff of S/N$>$4.5. All of the candidate lines are spatially unresolved, with implied radii of $<0.5"$ ($<$3\,kpc). Most of them are located away from the edges of individual spectral windows; the latter could lead to spurious sources due to increased noise. None of the candidates are detected in the 1\,mm continuum, which is consistent with the recent finding by \citet{capak15} that [CII] lines are more easily detected in z$\sim$5 Lyman--break galaxies than the underlying dust continuum. When stacking all high--redshift drop--out galaxies, we derive a 3$\sigma$ upper limit for the continuum emission of 6\,$\mu$Jy\,beam$^{-1}$. When stacking only at the location of the [CII] line candidates, we find a tentative detection of the dust continuum with a flux of $14\pm5\ \mu$Jy. This implies a dust--obscured star formation rate of 3\,M$_\odot$\,yr$^{-1}$. We compare the [CII] luminosities of our line candidates with the UV--based star formation rates and compare these with relations that have recently been discussed in the literature. These include galaxies and starbursts, as well as lower--metallicity dwarf/irregular galaxies. We find that the three highest--SFR objects have candidate [CII] lines with luminosities that are consistent with the low--redshift $L_{\rm [CII]}$ vs.\ UV-derived SFR relation. The other candidates have significantly higher [CII] luminosities. Given our {\it fidelity} analysis, we expect that $60\%$ of these sources to be spurious. A possible conclusion would be that some of the sources have elevated [CII] fluxes compared to expectations based on individual SFRs. Similar such claims have recently been reported in the literature \citep{maiolino16}. Future, deeper observations with ALMA will shed light on this issue. We note that confirming the objects of interest will request significantly less time with ALMA than the original survey, as no mosaic, nor frequency scans, would be required. Approved, deeper ALMA cycle~3 data of the same field will also improve the reliability of some of the candidates presented here. Based on the available information, we put first constraints on the [CII] luminosity function at $z\sim6-8$. Even if only one of our line candidates was real (a scenario greatly favoured by our statistical analysis) we find a source density that is consistent with the value derived by \citet{swinbank14} based on blindly detected [CII] emitters at higher luminosities at z$\sim$4.4. However these numbers are in conflict with a local (z=0) [CII] luminosity function derived by \citet{swinbank14} assuming a varying [CII]/IR luminosity ratio with IR luminosity. They are consistent though with a [CII] luminosity function that assumes a constant [CII]/IR$=0.002$ luminosity ratio \citep{swinbank14}. On the other hand, the high--redshift constraints so far appear to give significantly higher number densities than the recent models by \citet{popping16}. The observations presented here demonstrate that even in ALMA early science, [CII] luminosities can be reached that enable the studies of some of the faintest HST drop--out galaxies at $6.0<z<8.0$. Future, deeper and wider surveys with ALMA will be needed to improve the significance of the detections, and to improve the overall number statistics. With the fully completed ALMA now available, these goals now appear to be within reach. The full UDF appears to be the best field choice for such a survey, as the highest--redshift galaxy population is best--characterised in this field. In the near future, optical/NIR spectroscopy from JWST Guaranteed Time efforts will also provide accurate redshifts for the highest redshift galaxies in the UDF that would eventually enable pushing [CII] line studies to unprecedented depths through stacking.
16
7
1607.06772
We present a search for [C II] line and dust continuum emission from optical dropout galaxies at z &gt; 6 using ASPECS, our Atacama Large Millimeter submillimeter Array Spectroscopic Survey in the Hubble Ultra-deep Field (UDF). Our observations, which cover the frequency range of 212-272 GHz, encompass approximately the range of 6 &lt; z &lt; 8 for [C II] line emission and reach a limiting luminosity of L <SUB>[C II]</SUB> ∼ (1.6-2.5) × 10<SUP>8</SUP> L <SUB>⊙</SUB>. We identify 14 [C II] line emitting candidates in this redshift range with significances &gt;4.5σ, two of which correspond to blind detections with no optical counterparts. At this significance level, our statistical analysis shows that about 60% of our candidates are expected to be spurious. For one of our blindly selected [C II] line candidates, we tentatively detect the CO(6-5) line in our parallel 3 mm line scan. None of the line candidates are individually detected in the 1.2 mm continuum. A stack of all [C II] candidates results in a tentative detection with S <SUB>1.2 mm</SUB> = 14 ± 5 μJy. This implies a dust-obscured star-formation rate (SFR) of (3 ± 1) M <SUB>⊙</SUB> yr<SUP>-1</SUP>. We find that the two highest-SFR objects have candidate [C II] lines with luminosities that are consistent with the low-redshift L <SUB>[C II]</SUB> versus SFR relation. The other candidates have significantly higher [C II] luminosities than expected from their UV-based SFR. At the current sensitivity, it is unclear whether the majority of these sources are intrinsically bright [C II] emitters, or spurious sources. If only one of our line candidates was real (a scenario greatly favored by our statistical analysis), we find a source density for [C II] emitters at 6 &lt; z &lt; 8 that is significantly higher than predicted by current models and some extrapolations from galaxies in the local universe.
false
[ "C II", "SUB>⊙</SUB", "SFR relation", "II", "line emission", "spurious sources", "Array Spectroscopic Survey", "SFR", "UDF", "optical dropout galaxies", "candidates", "Atacama Large Millimeter", "lt", "±", "our Atacama Large Millimeter submillimeter Array Spectroscopic Survey", "z", "the Hubble Ultra-deep Field", "[C II", "II]</SUB", "galaxies" ]
13.299351
7.645282
132
12410081
[ "Harko, T.", "Mak, M. K." ]
2016Ap&SS.361..283H
[ "Exact power series solutions of the structure equations of the general relativistic isotropic fluid stars with linear barotropic and polytropic equations of state" ]
18
[ "Department of Physics, Babes-Bolyai University, Cluj-Napoca, Romania; Department of Mathematics, University College London, London, United Kingdom", "Departamento de Física, Facultad de Ciencias Naturales, Universidad de Atacama, Copiapó, Chile" ]
[ "2017Ap&SS.362..131K", "2017EPJC...77...97S", "2018EPJC...78..241S", "2018Prama..91...75S", "2018ZNatA..73..805M", "2019AnHP...20..813A", "2019PTEP.2019e3E02S", "2019Prama..92...63R", "2020BrJPh..50..725S", "2020EPJP..135..377I", "2020IJMPD..2950082A", "2020PDU....2800549A", "2020PDU....3000632A", "2021RoAJ...31..201M", "2023SciAf..2001696N", "2024IJTP...63...61K", "2024arXiv240205461S", "2024arXiv240516986B" ]
[ "astronomy", "physics" ]
5
[ "General relativistic fluid sphere", "Exact power series solutions", "Linear barotropic equation of state", "Polytropic equation of state", "General Relativity and Quantum Cosmology", "Astrophysics - Solar and Stellar Astrophysics", "Mathematical Physics" ]
[ "1870AmJS...50...57L", "1926ics..book.....E", "1930MNRAS..91....4M", "1930MNRAS..91...63F", "1939PhRv...55..364T", "1959PhRv..116.1027B", "1964ApJ...140..434T", "1974PhRvL..32..324R", "1980Ap&SS..73..227M", "1980esef.book.....K", "1988RvMP...60..297B", "1998CoPhC.115..395D", "1999MNRAS.303..466R", "1999PhLB..467...40S", "2000AnPhy.286..292N", "2000GReGr..32..919S", "2000MPLA...15.2153M", "2001MNRAS.328..839H", "2002AnP...514....3H", "2002ChJAA...2..248M", "2002IJMPD..11..207M", "2003GReGr..35.1435D", "2003RSPSA.459..393M", "2004ASSL..306.....H", "2004IJMPD..13..149M", "2004IJMPD..13.1005D", "2004IJMPD..13.1441P", "2004NewA....9..467N", "2005Prama..65..185M", "2007JCAP...06..025B", "2008A&A...483..673C", "2008PhRvD..78j4026O", "2009APh....31..128L", "2010JNMP...17..503B", "2010MNRAS.404L..50M", "2012JMP....53f2503M", "2012PhRvD..85f4045F", "2012PhRvD..86f4011C", "2013EPJC...73.2585M", "2013GReGr..45.1951T", "2013PhRvD..88h4022H", "2013PhRvD..88h4023A", "2015Ap&SS.356..309B", "2015EPJC...75..225M", "2015EPJC...75..442B", "2016CaJPh..94.1093R", "2016EPJC...76..106B", "2016PhRvD..94f4070B", "2018ZNatA..73..805M" ]
[ "10.1007/s10509-016-2875-0", "10.48550/arXiv.1607.06877" ]
1607
1607.06877_arXiv.txt
Karl Schwarzschild was the first scientist to find the exact solution of the Einstein's gravitational field equations describing the interior of a constant density compact astrophysical object in 1916 \citep{Sch1}. The search for exact solutions describing static neutral, charged, isotropic or anisotropic stellar type configurations has continuously attracted the interests of physicists and mathematicians. A wide range of analytical solutions of the gravitational field equations describing the interior structure of the static fluid spheres were found in the past 100 years (for reviews of the interior solutions of Einstein's gravitational field equations see \citep{0,1b,2b}). Unfortunately, among these many found solutions, there are very few exact interior solutions of the field equations satisfying the required general physical conditions. The criteria for physical acceptability of an interior solution can be formulated as follows \citep{1b}: 1) the solutions must be integrated from the regular origin of the stars. 2) the pressure and the energy density be positive definite at the origin of the stars. 3) the pressure vanishes at the surface of the stars. 4) the pressure and the energy density be monotonically decreasing to the surface of the stars for all radius. 5) causality requirement is that the speed of sound cannot be faster than the speed of light inside the stars. 6) the interior metric should be joined continuously with the exterior Schwarzschild metric. Note that in the field of static spherically symmetric fluid spheres, an important bound on the mass-radius ratio for stable general relativistic stars was obtained in \cite{Bu}, given by $2GM/c^{2}R\leq 8/9$, where $M$ is the mass of the star as measured by its external gravitational field, and $R$ is the boundary radius of the star. The Buchdahl bound was generalized to include the presence of the cosmological constant as well as higher dimensions and electromagnetic fields in \citep{MCC,Pi1,Pi2,Pi3}. In recent years, many exact solutions of the field equations describing the interior structure of the fluid stars have been found by assuming the existence of the anisotropic pressure \citep{B1,B2,B3,B4,A1,A3,A4,A2}. Since there are three independent field equations representing the stellar model, after adding the anisotropy parameter to the model, one has more mathematical freedom, and hence it is easier to solve the field equations analytically. However, it may be unphysical to assume the existence of anisotropic stresses. For instance, in a compact star, although the radial pressure vanishes at the surface of the star, one still could postulate the tangential pressure to exist. While the latter does not alter the spherical symmetry, it may create some streaming fluid motions \cite{3}. Thus, in order to obtain a realistic description of stellar interiors in the following we assume that the matter content of dense general relativistic can be described thermodynamically by the energy density $\rho \left( r\right) $ and the isotropic pressure $p\left( r\right) $. Therefore, from a mathematical point of view the isotropic stellar models are governed by the three field equations for four unknowns: the $tt$ and $rr$ components of the metric tensor $\exp \left[ \nu \left( r\right) \right] $ and $\exp \left[ \lambda \left( r\right) \right] $, the energy density $\rho \left( r\right) $% , and the pressure $p\left( r\right) $ respectively. Thus, the general relativistic stellar problem is an underdetermined one. In order to close the system of field equations an equation of state must be imposed. Very recently, the isotropic pressure equation was reformulated as a Riccati equation. By using the general integrability condition for the Riccati equation proposed in \cite{M1,M2}, an exact non-singular solution of the interior field equations for a fluid star expressed in the form of infinite series was obtained in \cite{00}. The astrophysical analysis indicates that this power series solution can be used as a realistic model for static general relativistic high density objects, for example neutron stars. In 1939, Tolman rewrote the isotropic pressure equation as the exact differential form involving the metric tensor components, subsequently leading him to obtain the eight analytical solutions of the field equations \citep{Tolman}. However, in order to ensure not to violate the causality condition, in the present paper, we do not follow Tolman's approach. Alternatively, we need one more constraint to close the system of the equations and to satisfy the causality requirement. Hence in the present paper we assume first that the matter energy density $\rho \left( r\right) $ and the thermodynamic pressure $p\left( r\right) $ obey the linear barotropic equation of state given by% \begin{equation} p\left( r\right) =\gamma \rho \left( r\right) c^{2}, \label{0} \end{equation}% where $\gamma $ is the arbitrary constant satisfying the inequality $0\leq \gamma \leq 1$. A static interior solution of the field equations in isotropic coordinates with the equation of state (\ref{0}) was presented in \cite{MMH}. The structure and the stability of relativistic stars with the equation of state (\ref{0}) were studied in \cite{PH}. An exact analytical solution describing the interior of a charged strange quark star satisfying the MIT bag model equation of state $3p=\rho c^{2}-4B$, where $B$ is a constant, was found in \cite{BV} under the assumption of spherical symmetry and the existence of a one-parameter group of conformal motions. Numerical solutions of Einstein's field equation describing static, spherically symmetric conglomerations of a photon gas, forming so-called photon stars, were obtained in \cite{Sch}. The solutions imply a back reaction of the metric on the energy density of the photon gas. In \cite{Glen} it was pointed out that a class of objects called Radiation Pressure Supported Stars (RPSS) may exist even in Newtonian gravity. Such objects can also exist in standard general relativity, and they are called "Relativistic Radiation Pressure Supported Stars" (RRPSS). The formation of RRPSSs can take place during the continued gravitational collapse. Irrespective of the details of the contraction process, the trapped radiation flux should attain the corresponding Eddington value at sufficiently large $z>>1$. On the basis of Einstein's theory of relativity, the principle of causality, and Le Chatelier's principle, in \cite{Ruf} it was established that the maximum mass of the equilibrium configuration of a neutron star cannot be larger than $3.2M_{\odot}$. To obtain this result it was assumed that for high densities the equation of state of matter is given by $p=\rho c^2$. The absolute maximum mass of a neutron star provides a decisive method of observationally distinguishing neutron stars from black holes. There is a long history in the context of physics and astrophysics for the study of the polytropic equation of state, defined as \citep{Hor} \begin{equation} p\left( r\right) =K\rho ^{\Gamma }\left( r\right) . \end{equation}% Here $K$ is the polytropic constant, and the adiabatic index $\Gamma $ is defined as $\Gamma =1+1/n$, where $n$ is the polytropic index. Using the polytropic equation of state, the physicists have investigated the properties of the astrophysical objects in Newtonian gravity. Note that $K$ is fixed in the degenerate system for instance a white dwarf or a neutron star and free in a non-degenerate system. The hydrostatic equilibrium structure of a polytropic star is governed for spherical symmetry by the Lane-Emden equation \citep{Hor} \begin{equation} \frac{1}{x^{2}}\frac{d}{dx}\left( x^{2}\frac{dy}{dx}\right) +y^{n}=0, \label{F1} \end{equation}% where the dimensionless variables $y$ and $x$ are defined as \begin{equation} x^{2}=\frac{4\pi G\rho _{c}^{\frac{n-1}{n}}}{\left( 1+n\right) K}r^{2},y^{n}=% \frac{\rho }{\rho _{c}}, \end{equation}% respectively where $\rho _{c}$ and $r$ are the central density and the radius of the star, respectively, $G$ is the Newtonian gravitational constant, and $y$ is the dimensionless gravitational potential. The Lane-Emden Eq. (\ref{F1}) was first introduced by \cite{L} and later studied by \cite{E}, \cite{Fow} and \cite{Milne}, respectively. In order to ensure the regularity of the solution at the center of the sphere, Eq. (\ref{F1}) must be solved with the initial conditions given by \begin{equation} y\left( 0\right) =1,\left( \frac{dy}{dx}\right) _{x=0}=0. \end{equation}% It is well-known that the exact analytical solutions of Eq. (\ref{F1}) can only be obtained for $n=0,1,5$ \citep{Hor,1s}. However, not all solutions of Eq. (\ref{F1}) for $n=5$ were known until the year 2012, when all real solutions of Eq.~(\ref{F1}) for $n=5$ were obtained in terms of Jacobian and Weierstrass elliptic functions \citep{PM}. Two integrable classes of the Emden-Fowler equation of the type $z\rq{}\rq{}=A \chi ^{-\lambda-2}z^n$ for $\lambda=\frac{n-1}{2}$, and $\lambda=n+1$ were discussed in \cite{Mancas}. By using particular solutions of the Emden-Fowler equations both classes were reduced to the form $\ddot \nu +a \dot \nu +b(\nu-\nu^n)=0$, where $a$, $b$ depend only on $\lambda $, and $n$, respectively. For both cases the solutions can be represented in a closed parametric form, with some values of $n$ yielding Weierstrass elliptic solutions. It is generally accepted that the power series method is one of the powerful techniques in solving ordinary differential equations. Thus the Lane-Emden Eq.~(\ref{F1}% ) was solved by using a power series method in \citep{CA,IW,CH,MN}, where the convergence of the solutions was also studied. The polytropic equation of state has also been adopted to study the interior structure of the fluid stars in the framework of general relativity \citep{Tooper}. The solution of the gravitational field equations for relativistic static spherically symmetric stars in minimal dilatonic gravity using the polytropic equation of state was presented in \cite{PK}. The general formalism to model polytropic general relativistic stars with the anisotropic pressure was considered in \cite{LW}, and its stellar applications were also discussed. By solving the Tolman-Oppenheimer-Volkoff (TOV) equation, a class of compact stars made of a charged perfect fluid with the polytropic equation of state was analyzed in \cite{jj}. Exact solutions of the Einstein-Maxwell equation with the anisotropic pressure and the electromagnetic field in the presence of the polytropic equation of state were obtained in \cite{PS}. Charged polytropic stars, and a generalization of the Lane-Emden equation was investigated in \cite{RM}. Using the power series methodology, a new analytical solution of the TOV equation for polytropic stars was presented in \cite{MAS}. The divergence and the convergence of the power series solutions for the different values of the polytropic index $n$ were also discussed. The gravitational field equations for the static spherically symmetric perfect fluid models with the polytropic equation of state can be written as two complementary 3 dimensional regular systems of ordinary differential equations on compact state space. Due to the highly nonlinear structure of the systems, it is difficult to solve them exactly, and thus they were analyzed numerically and qualitatively using the theory of dynamical systems in \citep{UN,KCC}. The three-dimensional perfect fluid stars with the polytropic equation of state, matched to the exterior three-dimensional black hole geometry of Ba\~{n}ados, Teitelboim and Zanelli were considered in \cite{PTA}. A new class of exact solutions for a generic polytropic index was found, and analyzed. The structure of the relativistic polytropic stars and the stellar stability analysis embedded in a chameleon scalar field was discussed in \cite{VD}. In \cite{XY} a polytropic quark star model was suggested in order to establish a general framework in which theoretical quark star models could be tested by the astrophysical observations. Spherically symmetric static matter configurations with the polytropic equation of state for a class of $f\left( R\right) $ models in Palatini formalism were investigated in \cite{GO}, and it was shown that the surface singularities are not physical in the case of Planck scale modified Lagrangians. It is the purpose of the present paper to study the interior structure of the general relativistic fluid stars with the linear barotropic and the polytropic equations of state, and to obtain exact power series solutions of the corresponding equations. As a first step in our study we introduce the basic equation describing the interior mass profile of a relativistic star, and which we call the relativistic mass equation. This equation is obtained by eliminating the energy density between the mass continuity equation and the hydrostatic equilibrium equation. Despite its apparent mathematical complexity, the relativistic mass equation can be solved exactly for both linear barotropic and polytropic equations of state, by looking to its exact solutions as represented in the form of power series. In order to obtain closed form representations of the coefficients we use the Cauchy convolution of the power series. In this way we obtain the exact series solutions for relativistic spheres described by linear barotropic equations of state with arbitrary $\gamma $, and for the polytropic equation of state with arbitrary polytropic index $n$. The case $n=1$ is investigated independently, and the corresponding power series solution is also obtained. We compare the truncated power series solutions containing seven terms only with the exact numerical solution of the TOV and mass continuity equations. In all considered cases we find an excellent agreement between the power series solution, and the numerical one. The present paper is organized as follows. The gravitational field equations, their dimensionless formulation and the basic relativistic mass equation are presented in Section~\ref{sect1}. The definition of the Cauchy convolution for infinite power series is also introduced. The non-singular power series solution for fluid spheres described by a linear barotropic equation of state is presented in Section~\ref{sect2}. The comparison between the exact and numerical solutions are presented. The exact power series solutions for a general relativistic polytropic star with polytropic index $n=1$ are derived in Section~\ref{sect3}, and the comparison with the exact numerical solution is also performed. The case of the arbitrary polytropic index $n$ is considered in Section~\ref{sect4}. The power series solutions are compared with the exact numerical solutions for the cases $% n=1/2$, $n=1/5$ and $n=3$, respectively. We discuss our results and conclude our paper in Section~\ref{sect5}. The first seven coefficients of the power series solution of the relativistic mass equation for arbitrary polytropic index $n$ are presented in Appendix~\ref{app}.
\label{sect5} In the present paper we have obtained exact power series solutions of the mass continuity and hydrostatic equilibrium equation describing the structure of general relativistic stars. In order to obtain the solutions we have formulated the second order differential equation describing the mass profile of the stars. The relativistic mass equation admits exact, convergent and non-singular, power series solutions for both the linear barotropic and polytropic equations of state. We have obtained the power series solutions for arbitrary values of $\gamma $ for the linear barotropic equation of state $p=\gamma \rho c^{2}$, and for the general case of the arbitrary polytropic index $n$. We have compared in detail our exact results with the results obtained by numerically integrating the gravitational field equations, by considering the cases $\gamma =1/3$, $\gamma =1$, and $% n=1,1/2,1/5,3$, respectively. By truncating our power series to only seven terms we can basically reproduce the numerical results for the mass and density distribution of the general relativistic stars described by linear barotropic and polytropic equations of state. The power series solution are non-singular at the center of the star, and they can be extended up to the vacuum boundary/surface of the dense matter distribution. Due to the adopted equations of state the physical requirements for the acceptability of the solutions are automatically satisfied. Thus, the speed of sound $c_{s}=\sqrt{% \partial p/\partial \rho }=\sqrt{\gamma }c$ is a constant inside the star, and for $\gamma \in \lbrack 0,1]$ satisfies the constraint $c_{s}\leq c$. For the polytropic stars we obtain \be c_{s}=\sqrt{k(n+1)/n}\rho ^{1/2n}=c_{s}\left( \rho _{c}\right) \epsilon ^{1/2n}, \ee where we have denoted \be c_{s}\left( \rho _{c}\right) =\rho _{c}^{1/2n}\sqrt{k(n+1)/n}. \ee Using the relation $\epsilon \left( \eta \right) =\sum_{i=1}^{\infty }\left( 2i+1\right) c_{2i+1}\eta ^{2i-2}$ then we find \be c_{s}=c_{s}\left( \rho _{c}\right) \Bigg[ \sum_{i=1}^{\infty }{\left( 2i+1\right) c_{2i+1}\eta ^{2i-2}}\Bigg] ^{ 1/2n}\leq c. \ee The power series solutions of the Newtonian Lane-Emden Eq.~(\ref{F1}) have been intensively studied in the astrophysical and mathematical literature \citep{CA,IW,CH,MN}. The series solutions are represented as $\theta =\sum _{k=0}^{\infty}{a_k\xi ^{2k}}$ and $\theta ^n= \sum _{k=0}^{\infty}{b_k\xi ^{2k}}$, respectively, with $a_0=b_0=1$ \citep{MN}. One can define the radius of convergence of these series as the distance from $\xi = 0$ to the closest singularity of $\theta (\xi)$ in the complex $\xi $-plane. Non-linear ordinary differential equations, such as the Lane–Emden equation for $n>1$, can have two kinds of singularities, fixed and movable \citep{MN}. The Lane–Emden equation for polytropic index $n>1$ and its $n\rightarrow \infty$ limit, corresponding to the limit of the isothermal sphere equation, are singular at some negative value of the radius squared \citep{MN}. It is this singularity that prevents real power series solutions about the center to converge to the outer surface once the condition $n>1.9121$ is satisfied. However, as shown in \cite{MN}, an Euler transformation gives power series that do converge up to the outer radius. Moreover, the Euler-transformed series converge significantly faster than the series obtained in \cite{IW}, which are limited to finite radii whenever $n>5$ by a complex conjugate pair of singularities. Series solutions for polytropic stars by using the Euler transform were constructed in \cite{MN}, so that longer than 60-term series are needed for the outer regions of $n>3$ polytropic Newtonian stars, while 120-term and 300-term series are needed to obtain the function $\theta (\xi)$ to seven decimal place accuracy all the way from the center to the surface of the compact object for $n=3.5$ and $% n=4 $, respectively \citep{MN}. In this context we would like to point out that the power series solutions of the relativistic mass equation can be extended to the boundary of the considered stars, and only seven terms are required to reproduce the numerical solutions with a high precision. Although the numerical solutions of the structure equations of spherically symmetric static general relativistic stars can be obtained numerically in a very efficient, simple and accurate way, we must point out that power series represent one of the most powerful methods of mathematical analysis. The use of power series is no less convenient than the use of elementary functions, especially when solutions of differential equations are to be studied numerically. In the case of the approach based on the relativistic mass equation an important advantage of a power series solution is that it gives the value of the mass and energy density inside the star as a recurrent power series in the radial coordinate $r$, since the dimensionless variable $% \eta =r/R$. Consequently, we can predict the physical and geometrical parameters of the star at any radius directly. Moreover, power series analytical solutions describing the interior of compact general relativistic objects usually offers deeper insights into their physical and geometrical properties, thus offering the possibility of a better understanding of the structure of dense stars.
16
7
1607.06877
Obtaining exact solutions of the spherically symmetric general relativistic gravitational field equations describing the interior structure of an isotropic fluid sphere is a long standing problem in theoretical and mathematical physics. The usual approach to this problem consists mainly in the numerical investigation of the Tolman-Oppenheimer-Volkoff and of the mass continuity equations, which describes the hydrostatic stability of the dense stars. In the present paper we introduce an alternative approach for the study of the relativistic fluid sphere, based on the relativistic mass equation, obtained by eliminating the energy density in the Tolman-Oppenheimer-Volkoff equation. Despite its apparent complexity, the relativistic mass equation can be solved exactly by using a power series representation for the mass, and the Cauchy convolution for infinite power series. We obtain exact series solutions for general relativistic dense astrophysical objects described by the linear barotropic and the polytropic equations of state, respectively. For the polytropic case we obtain the exact power series solution corresponding to arbitrary values of the polytropic index n. The explicit form of the solution is presented for the polytropic index n=1, and for the indexes n=1/2 and n=1/5, respectively. The case of n=3 is also considered. In each case the exact power series solution is compared with the exact numerical solutions, which are reproduced by the power series solutions truncated to seven terms only. The power series representations of the geometric and physical properties of the linear barotropic and polytropic stars are also obtained.
false
[ "exact series solutions", "infinite power series", "exact solutions", "general relativistic dense astrophysical objects", "the exact power series solution", "the power series solutions", "the relativistic mass equation", "n=1/5", "The power series representations", "a power series representation", "the polytropic equations", "the mass continuity equations", "the spherically symmetric general relativistic gravitational field equations", "the exact numerical solutions", "arbitrary values", "the polytropic index", "the linear barotropic and polytropic stars", "the relativistic fluid sphere", "theoretical and mathematical physics", "state" ]
8.942719
3.890101
28
901417
[ "Barnes, David J.", "Kay, Scott T.", "Henson, Monique A.", "McCarthy, Ian G.", "Schaye, Joop", "Jenkins, Adrian" ]
2017MNRAS.465..213B
[ "The redshift evolution of massive galaxy clusters in the MACSIS simulations" ]
105
[ "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, UK", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, UK", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, UK", "Astrophysics Research Institute, Liverpool John Moores University, 146 Brownlow Hill, Liverpool L3 5RF, UK", "Leiden Observatory, Leiden University, PO Box 9513, NL-2300 RA Leiden, the Netherlands", "Institute for Computational Cosmology, Department of Physics, University of Durham, South Road, Durham DH1 3LE, UK" ]
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[ "astronomy" ]
27
[ "hydrodynamics", "methods: numerical", "galaxies: clusters: general", "galaxies: clusters: intracluster medium", "galaxies: evolution", "X-rays: galaxies: clusters", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1093/mnras/stw2722", "10.48550/arXiv.1607.04569" ]
1607
1607.04569_arXiv.txt
\label{sec:intro} Galaxy clusters form from large primordial density fluctuations that have collapsed and virialised by the present epoch, with more massive clusters forming from larger and rarer fluctuations. This makes them especially sensitive to fundamental cosmological parameters, such as the matter density, the amplitude of the matter power spectrum and the equation of state of dark energy \citep[see][]{Voit2005,AllenEvrardMantz2011,KravtsovBorgani2012,Weinberg2013}. The observable properties of a galaxy cluster result from a non-trivial interplay between gravitational collapse and astrophysical processes. The diverse range of formation histories of the cluster population leads to scatter in the observable-mass scaling relations and, as surveys select clusters based on an observable, this can lead to a biased sample of clusters, resulting in systematics when using them as a cosmological probe \citep[e.g.][]{Mantz2010}. Many previous studies have shown that the relationship between a cluster observable, such as its temperature or X-ray luminosity, and a quantity of interest for cosmology, e.g. its mass, has a smaller scatter for more massive, dynamically relaxed objects \citep{EkeNavarroFrenk1998,Kay2004,Crain2007,NagaiVikhlininKravtsov2007a,Planelles2013}. Therefore, the fundamental requirement when probing cosmological parameters with galaxy clusters is a sample of relaxed, massive clusters with well calibrated mass-observable scaling relations. However, galaxy clusters are rare objects, becoming increasingly rare with increasing mass, and to observe a sample large enough to be representative of the underlying population requires a survey with significant size and depth. Currently ongoing and impending observational campaigns, such as the Dark Energy Survey \citep{DEScol2005}, \textit{eRosita} \citep{Merloni2012}, \textit{Euclid} \citep{Laureijs2011}, \textsc{SPT-3G} \citep{Benson2014} and Advanced ACTpol \citep{Henderson2016}, will be the first to have sufficient volume to yield significant samples of massive clusters. Due to their rarity, the majority of these massive clusters will be at high redshift and it is therefore critical to understand how the cluster observables and their associated scatter evolve. Additionally, the most massive clusters will be the brightest and easiest to detect objects at high redshift, making it vital to understand the selection function of the chosen cluster observable and whether the most massive clusters are representative of the underlying cluster population. Theoretical modelling of the formation of clusters and their observable properties is required to understand these issues and to further clusters as probes of cosmology. Due to the range of scales involved in cluster formation, the need to incorporate astrophysical processes and to self-consistently predict observable properties, cosmological hydrodynamical simulations are the only viable option. Recent progress in the modelling of large-scale structure formation has been driven mainly by the inclusion of supermassive black holes and their associated Active Galactic Nucleus (AGN) feedback, which has been shown to be critical for reproducing many cluster properties \citep{Bhattacharya2008,PuchweinSijackiSpringel2008,McCarthy2010,Fabjan2010}. A number of independent simulations are now able to produce realistic clusters that simultaneously reproduce many cluster properties in good agreement with the observations \citep{LeBrun2014,Pike2014,Planelles2014}. Results from the recent BAryons and HAloes of MAssive Systems (BAHAMAS) simulations \citep{McCarthy2016} have shown that by calibrating the subgrid model for feedback to match a small number of key observables, in this case the global galaxy stellar mass function and the gas fraction of clusters, simulations of large-scale structure are now able to reproduce many observed scaling relations and their associated scatter over two decades in halo mass. However, full gas physics simulations of large-scale structure formation, with sufficient resolution, are still computationally expensive. This has limited previous studies to either small samples with $<50$ objects or to volumes of $596\,\mathrm{Mpc}$, all of which are too small to contain the representative sample of massive clusters that is required for cosmological studies above $z=0$. This paper introduces the Virgo consortium's MACSIS project, a sample of $390$ massive clusters selected from a large volume dark matter simulation and resimulated with full gas physics to enable self-consistent observable predictions. The simulations extend the BAHAMAS simulations to the most massive clusters expected to form in a $\Lambda\rm{CDM}$ cosmology. In this paper we study the cluster scaling relations and their evolution. We combine the MACSIS and BAHAMAS simulations to produce a sample that spans the complete mass range and that can be studied to high redshift, using the progenitors of the MACSIS sample. We also select the hottest clusters from the combined sample and a relaxed subset of them to examine the impact of such selections on the scaling relations and their evolution. We then study the gas profiles to further understand the differences between the samples. This paper is organised as follows. In Section \ref{sec:MACsamp} we introduce the MACSIS sample and discuss the parent dark matter simulation from which the sample was selected, the selection criteria used, the model used to resimulate the haloes, how we produced the observable quantities and the three samples we use in this work. In Section \ref{sec:screlations} we investigate how the scaling relations evolve and how this evolution changes when a hot cluster sample or relaxed, hot cluster sample is selected. We then study the hot gas profiles to understand the differences in the evolution of the relations for the different samples in Section \ref{sec:gasprofs}. Finally, in Section \ref{sec:sad} we discuss our results and summarise our main findings.
\label{sec:sad} In this work we have presented the MACSIS clusters, a sample of 390 zoomed simulations of the most massive and rarest clusters run with the state-of-the-art, calibrated baryonic physics model from the BAHAMAS project \citep{McCarthy2016} that yields realistic clusters. Such massive clusters are absent from the BAHAMAS simulation volumes of $596\,\mathrm{Mpc}$ as the simulated volume is too small. After introducing the selection of the sample from the parent $3.2\,\mathrm{Gpc}$ volume simulated with the \textit{Planck} 2013 cosmology, and demonstrating the agreement of the properties of our massive cluster sample with the properties of observed massive clusters, we examined the evolution of the cluster scaling relations and the evolution of the cluster gas profiles. By combining the MACSIS sample with the clusters in the BAHAMAS volume, we were able to examine the cluster scaling relations over the full observed mass range for the first time. Additionally, the MACSIS clusters enabled the study of the evolution of the cluster scaling relations to unprecedentedly high redshifts. Finally, the MACSIS sample enabled clusters to be selected in ways which mimic a cosmological study, such as selecting the hottest clusters, to examine if the scaling relations of such objects evolve differently from the underlying cluster population. Our main results are: \begin{itemize} \item As shown in Fig. \ref{fig:observations}, the MACSIS simulations yield realistic massive clusters at low redshift and their progenitors are in good agreement with the limited observational data that is available at high redshift (i.e. $z=1$). \item Scaling relations for the combined sample that spans the full observed cluster mass range show significant deviations from the simple self-similar theory (see Figs. \ref{fig:MgMsr}-\ref{fig:LxTxsr}). Both the slope of the relations and the redshift evolution of the normalization are significantly affected by non-gravitational physics. The low redshift relations are in good agreement with observations and with most previous simulation work. \item The main drivers of non-self-similar evolution are AGN feedback, non-thermal pressure support and a mild mass dependence of the spectroscopic temperature bias. Shallower potentials of clusters that are less massive or form at lower redshifts allows feedback from AGN to eject more gas. Non-thermal pressure lowers a cluster's temperature for a given potential and is more important in more massive clusters that have had less time to thermalise. We found that the spectroscopic temperature bias increases for the most massive clusters. \item With the exception of the luminosity-temperature relation, we found the scatter about the best-fit scaling relations is insensitive to mass and redshift for all of the cluster samples. \item Selecting a hot cluster sample, i.e. core-excised spectroscopic temperatures $k_{\rm{B}}T^{\mathrm{X,ce}}_{500\mathrm{spec}}\geq5\,\rm{keV}$, significantly alters the scaling relations and their evolution. Excluding the spectroscopic temperature-total mass relation, we find that the scaling relations of the hot cluster sample evolve in a much more self-similar manner. After accounting for the expected self-similar evolution with redshift, we find that the normalizations are consistent with no evolution. The slopes of the best-fit relations at each redshift are also broadly consistent with the slopes predicted by self-similar theory. However, the spectroscopic temperature-total mass relation of the hot sample deviates further from self-similarity than the combined sample. Selecting hot clusters removes the less massive clusters from the sample, so the hot sample is dynamically younger than the combined sample as more massive clusters form later in the hierarchical merger scenario. This increases the average level of non-thermal support in the hot sample, leading to a flatter spectroscopic temperature-total mass relation. Additionally, the spectroscopic temperature bias flattens the relation for the most massive clusters and this has a larger impact in a sample of only hot clusters. \item Selecting a relaxed subset of hot clusters, where the most dynamically disturbed objects are removed, leads to a small reduction in the scatter for most scaling relations. Removing the most disturbed objects also leads to a reduction in the level of non-thermal support in the sample compared to the complete hot sample. This leads to steeper slope of the spectroscopic temperature-total mass relation compared to the hot sample and a value that is closer to the self-similar prediction. \item The median hot gas profiles of the combined sample in general shows good agreement with observed radial profiles. The low redshift data is in very good agreement, while the data at $z=1$ shows reasonable agreement with the relaxed hot sample. \item Comparison of the hot gas profiles at $z=0$ and $z=1$ show evolution different from self-similar prediction (see Figs. \ref{fig:gas_prof}-\ref{fig:enty_prof}). The combined sample shows a decreasing density profile with decreasing redshift, suggesting the impact of AGN feedback. Selecting a sample of hot clusters produces a median density profile that evolves in much more self-similar manner. The combined and hot samples have a median temperature profile that increases with decreasing redshift. This is likely driven by decreasing importance of non-thermal pressure support with decreasing redshift. Selecting relaxed hot cluster sample produces a median profile that evolves in better agreement with the self-similar prediction. \end{itemize} MACSIS enables the study of the observable properties of the most massive and rarest galaxy clusters. We have demonstrated that their progenitors provide a good match to the currently limited observational data at high redshift and that their observable properties evolve in a significantly more self-similar manner than for lower-mass and less-relaxed clusters. We have shown how the selection function can impact the derived scaling relations and radial profiles. The size of the parent simulation enables the creation of synthetic lightcones with an area comparable to currently ongoing surveys. This will allow the impact of selection biases to be fully examined and the covariance of observable properties to be studied. Another route for future work is to improve our understanding of structure in the ICM, as the limited resolution and traditional SPH scheme used in this work limits our ability resolve structures and understand their impact on observable properties.
16
7
1607.04569
We present the MAssive ClusterS and Intercluster Structures (MACSIS) project, a suite of 390 clusters simulated with baryonic physics that yields realistic massive galaxy clusters capable of matching a wide range of observed properties. MACSIS extends the recent BAryons and HAloes of MAssive Systems simulation to higher masses, enabling robust predictions for the redshift evolution of cluster properties and an assessment of the effect of selecting only the hottest systems. We study the observable-mass scaling relations and the X-ray luminosity-temperature relation over the complete observed cluster mass range. As expected, we find that the slope of these scaling relations and the evolution of their normalization with redshift depart significantly from the self-similar predictions. However, for a sample of hot clusters with core-excised temperatures k<SUB>B</SUB>T ≥ 5 keV, the normalization and the slope of the observable-mass relations and their evolution are significantly closer to self-similar. The exception is the temperature-mass relation, for which the increased importance of non-thermal pressure support and biased X-ray temperatures leads to a greater departure from self-similarity in the hottest systems. As a consequence, these also affect the slope and evolution of the normalization in the luminosity-temperature relation. The median hot gas profiles show good agreement with observational data at z = 0 and z = 1, with their evolution again departing significantly from the self-similar prediction. However, selecting a hot sample of clusters yields profiles that evolve significantly closer to the self-similar prediction. In conclusion, our results show that understanding the selection function is vital for robust calibration of cluster properties with mass and redshift.
false
[ "cluster properties", "hot clusters", "robust predictions", "clusters", "realistic massive galaxy clusters", "observed properties", "redshift depart", "biased X-ray temperatures", "evolution", "redshift", "the complete observed cluster mass range", "mass", "non-thermal pressure support", "the redshift evolution", "the self-similar prediction", "the self-similar predictions", "the X-ray luminosity-temperature relation", "robust calibration", "baryonic physics", "MAssive Systems simulation" ]
13.036512
4.727173
163
12409466
[ "Zahid, H. Jabran", "Geller, Margaret J.", "Fabricant, Daniel G.", "Hwang, Ho Seong" ]
2016ApJ...832..203Z
[ "The Scaling of Stellar Mass and Central Stellar Velocity Dispersion for Quiescent Galaxies at z&lt;0.7" ]
74
[ "Smithsonian Astrophysical Observatory, Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Smithsonian Astrophysical Observatory, Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Smithsonian Astrophysical Observatory, Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "School of Physics, Korea Institute for Advanced Study, 85 Hoegiro, Dongdaemun-gu, Seoul 02455, Korea" ]
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[ "astronomy" ]
13
[ "dark matter", "galaxies: evolution", "galaxies: formation", "galaxies: high-redshift", "galaxies: structure", "Astrophysics - Astrophysics of Galaxies" ]
[ "1972CoASP...4..173S", "1976ApJ...204..668F", "1980AJ.....85..801S", "1987ApJ...313...42D", "1987ApJ...313...59D", "1995MNRAS.276.1341J", "1998ApJ...508L..43F", "1999ApJ...527...54B", "1999MNRAS.310..540A", "2000AJ....120.1579Y", "2000ApJ...533..682C", "2002AJ....123..485S", "2002AJ....124.1810S", "2002SPIE.4836...73W", "2003A&A...407..423M", "2003MNRAS.341...33K", "2003MNRAS.344.1000B", "2003PASP..115..763C", "2004A&A...425..881L", "2004MNRAS.351.1151B", "2004MNRAS.353..189V", "2004PASP..116..138C", "2005AJ....129...61B", "2005AJ....129.2562B", "2005ApJ...635L.125G", "2005PASP..117.1411F", "2006A&A...457..841I", "2006ApJ...647..201C", "2006MNRAS.366.1126C", "2006MNRAS.371..703S", "2007AJ....133..734B", "2008ApJ...672..122E", "2008ApJ...684..248B", "2008Natur.455.1082D", "2008PASP..120.1222F", "2009A&A...501.1269K", "2009ASPC..411..251M", "2009ApJ...699..486C", "2009MNRAS.394.1978H", "2009MNRAS.396.1171H", "2010AJ....139.1628D", "2010AJ....139.1857W", "2010ApJ...709.1195T", "2010MNRAS.404.1111G", "2010Natur.468..940V", "2011A&A...534A..61N", "2011MNRAS.410..210M", "2011MNRAS.412L...6B", "2011MNRAS.418.2785M", "2012AJ....143...90S", "2012ApJ...751L..44W", "2012Natur.484..485C", "2012arXiv1201.1913W", "2013A&A...549A...7V", "2013A&A...556A..55I", "2013AJ....146...32S", "2013ApJ...762L...7L", "2013ApJ...766...33L", "2013ApJ...767...50M", "2013ApJ...770...57B", "2013ApJ...777...18M", "2013ApJ...777L..10B", "2013ApJ...779L..21B", "2013MNRAS.431.1383T", "2013MNRAS.432.1862C", "2013MNRAS.433.3017L", "2013PASP..125.1362F", "2014ApJ...783..117B", "2014ApJ...784..162P", "2014ApJ...791..130Z", "2014ApJS..213...35G", "2014MNRAS.439.2505B", "2014MNRAS.440..610P", "2014MNRAS.443L..69S", "2014MNRAS.444.1518V", "2015ApJ...800..124B", "2015ApJS..219...12A", "2015MNRAS.446..521S", "2015MNRAS.449.3441S", "2015MNRAS.450.1349K", "2015MNRAS.450.3696M", "2015MNRAS.454.2770T", "2016A&A...594A..13P", "2016ARA&A..54..597C", "2016ApJ...818..173H", "2016ApJ...819...63R", "2016ApJ...821..101Z", "2016ApJS..224...11G", "2016MNRAS.456.3265M" ]
[ "10.3847/0004-637X/832/2/203", "10.48550/arXiv.1607.04275" ]
1607
1607.04275_arXiv.txt
In the standard model of cosmology the vast majority of matter in the universe ($\sim84\%$) is dark \citep{Planck2015}. The large scale structure of the universe develops as gravity acts on small density fluctuations in the virtually uniform dark matter dominated initial matter distribution. Gravity forms and shapes dark matter into nearly spherical halos. Baryons accrete onto dark matter halos and subsequently cool and condense to form galaxies. Dark matter halos are more extended and substantially more massive than the galaxies which form and evolve at their centers. Within the hierarchical formation paradigm, large scale structure formation and galaxy evolution are primarily driven by the accretion of dark matter and by halo mergers. Dark matter can not be directly observed and thus a central issue for cosmology is what observable property is the best proxy for connecting galaxies$-$the visible tracers of the matter distribution$-$to the dark matter distribution. Within the broader cosmological context, connections among the observed properties of galaxies should elucidate the physical processes governing galaxy formation and evolution. Statistical analyses of the galaxy population have established that the principal observable galaxy properties are all correlated, though the fundamental parameter in these correlations remains uncertain \citep{Disney2008}. Several recent studies suggest that either stellar mass or stellar velocity dispersion is the fundamental parameter characterizing galaxies and their dark matter halos \citep{More2011, Wake2012a, Wake2012b, Li2013, vanUitert2013, Bogdan2015}. The stellar mass and velocity dispersion are governed by different physical processes, but both are intimately related to properties of the dark matter halo. Stellar mass and velocity dispersion are strongly correlated making it difficult to determine which of these two parameters is fundamental. The stellar mass is an integral over the star formation history and the end product of the complex baryonic processes governing galaxy formation and evolution. In contrast, the velocity dispersion is a measure of the stellar kinematics and is directly related to the gravitational potential of the system. The correlation between luminosity and velocity dispersion in elliptical galaxies is well established. Based on observations of 25 galaxies, \citet{Faber1976} find a power law relation between velocity dispersion and luminosity and conclude that the total mass is the most fundamental property of elliptical galaxies. The stellar mass-to-light ratio for elliptical galaxies does not vary significantly; a relation between stellar mass and velocity dispersion directly follows from the \citet{Faber1976} result. Many subsequent studies based on larger samples and/or spatially resolved spectroscopy confirm a power law relation between stellar mass and velocity dispersion {over most of the stellar mass range explored in these studies} \citep[e.g.,][]{Hyde2009a, Cappellari2013b, Cappellari2016}. The redshift evolution of the relation between stellar mass and velocity dispersion and the scatter around the relation provide further constraints for models of galaxies. \citet{Belli2014} show that the relation between stellar mass and velocity dispersion at $0.9<z<1.6$ is offset from the local relation. \citet{Belli2014} attribute this offset to the smaller sizes of galaxies at higher redshift. This interpretation is consistent with the fact that the intrinsic scatter in the relation between luminosity and stellar mass is strongly correlated with size$-$the so-called fundamental plane \citep{Djorgovski1987, Dressler1987}; this relation also exists when luminosity is replaced by stellar mass \citep{Hyde2009b}. The stellar mass fundamental plane does not appear to evolve strongly for $z<0.6$ \citep{Zahid2016} but may at higher redshifts \citep{Bezanson2013}. Thus, the smaller sizes of galaxies at high redshift may explain their larger velocity dispersions at a fixed stellar mass. \citet{Shu2012} analyze the velocity dispersion distribution of a large sample of luminous red galaxies observed as part of the Baryon Oscillation Spectroscopic Survey (BOSS). Due to the limiting signal-to-noise of their observations, \citet{Shu2012} measure the intrinsic scatter by a ``Bayesian stacking" technique rather than from examination of the distribution of individual galaxy velocity dispersions. They report that the intrinsic scatter in velocity dispersions at a fixed stellar mass increases as a function of redshift, though the evolution is small for $z\lesssim0.6$ \citep[see Figure 11 in][]{Shu2012}. They conclude that the increased scatter indicates greater diversity in the galaxy population at early times. {In a more recent analysis of the BOSS sample, \citet{Montero-Dorta2016} find that the slope and scatter in the relation between luminosity and velocity dispersion for high-mass red sequence galaxies does not evolve significantly between $0.5<z<0.7$. Thus, studies of large samples of red galaxies from BOSS indicate little evolution in the relation between velocity dispersions and stellar mass or luminosity.} Here we analyze a sample of 4585 galaxies with individual stellar mass and velocity dispersion measurements at $z<0.7$ to examine the stellar mass-velocity dispersion ($M_\ast \sigma$) relation and its dependence on redshift. The intermediate redshift range we probe connects SDSS observations to the higher redshift observations of \citet{Belli2014} and our analysis of the scatter around the $M_\ast \sigma$ relation is complementary to \citet{Shu2012}; we measure a velocity dispersion for each galaxy. In Section 2 we describe the sample and present the $M_\ast \sigma$ relation in Section 3. We analyze the scatter in the $M_\ast \sigma$ relation in Section 4 and highlight important systematic issues in Section 5. We discuss the results in Section 6 and summarize and conclude in Section 7. We adopt the standard cosmology $(H_{0}, \Omega_{m}, \Omega_{\Lambda}) = (70$ km s$^{-1}$ Mpc$^{-1}$, 0.3, 0.7) throughout.
We examine the relation between stellar mass and central stellar velocity dispersion$-$the $M_\ast \sigma$ relation$-$for massive quiescent galaxies at $z<0.7$ using data from SDSS and SHELS \citep[cf.][]{Montero-Dorta2016}. We cross-calibrate measurements of stellar mass and velocity dispersion for these two surveys. The $M_\ast \sigma$ relation and its scatter are both independent of redshift. The central stellar velocity dispersion and stellar mass may provide comparable descriptions of galaxies at $z<0.7$. However, the scatter in the relation between stellar mass and velocity dispersion correlates with other galaxy physical properties. There is also tentative evidence that the zero point of the relation evolves at $z>1$. In detail, the velocity dispersion and stellar mass are not equivalent descriptors. Determining which of these properties is more closely related to the dark matter halo has important implications for understanding galaxy evolution and cosmology. The scaling between stellar mass and velocity dispersion may provide important clues for connecting early-type galaxies to a broader class of objects. The $M_\ast \sigma$ relation goes as $\sigma \propto M_\ast^{0.3}$. This scaling is expected for virialized dark matter halos, i.e. $\sigma_{DM} \propto M_{200}^{0.33}$ where $\sigma_{DM}$ is the dark matter halo velocity dispersion and $M_{200}$ is the total virial mass \citep{Evrard2008}. This similarity is somewhat surprising because there is no a priori expectation that the stellar mass and the stellar velocity dispersion should scale in the same manner as the dark matter halo mass and the dark matter halo velocity dispersion. We infer the dark matter halo mass of galaxies using a standard stellar-to-halo-mass conversion \citep{Behroozi2013a}. The relation between the total galaxy mass (stellar + dark matter) and the stellar velocity dispersion is surprisingly consistent with an extrapolation of the relation between the total mass of a galaxy cluster and its velocity dispersion \citep{Rines2016}. This result suggests that the stellar velocity dispersion may be directly proportional to the dark matter halo velocity dispersion. Our results set the stage for more detailed explorations of the connection between stellar velocity dispersions and the properties of dark matter halos. The velocity dispersion may be a more fundamental observable than stellar mass and may provide a more robust means for connecting dark matter halos in N-body simulations with observed galaxies. A more direct connection between observations and theory could be based on a targeted analysis of large cosmological hydrodynamic simulations like the Illustris, EAGLE and MassiveBlack-II projects \citep{Vogelsberger2014a, Schaye2015, Khandai2015}. The luminosity-weighted two dimensional line-of-sight central velocity dispersion of stellar particles within a circular aperture (i.e. the projected velocity dispersion measured within a cylinder in analogy with the observations) could be compared with the appropriately measured velocity dispersion of particles characterizing the dark matter halo. The relationship between these two velocity dispersions derived from the simulations could provide a guide for better interpretation of the observations. \begin{deluxetable*}{llcccc} \tablewidth{500pt} \tablecaption{Fit Parameters} \tablehead{\colhead{Sample} & \colhead{N} &\colhead{log($M_b/M_\odot$)} &\colhead{log($\sigma_b$) [km s$^{-1}$]} &\colhead{$\alpha_1$} & \colhead{$\alpha_2$} } \startdata \cutinhead{The full sample (Figure \ref{fig:fit}) } SDSS & 371884 & $10.26 \pm 0.01$ & $2.073 \pm 0.003$ & $ 0.403 \pm 0.004 $& $0.293 \pm 0.001$ \\ SHELS & 4585 & 10.26 & $2.071 \pm 0.004$ & & $0.281\pm0.005$ \\ \cutinhead{The full sample as a function of redshift (Figure \ref{fig:fitz})} SDSS ($M_\ast > 10^{10.26} M_\odot$) &325186 & 11 & $2.2969 \pm 0.0006$ && $0.299 \pm 0.001$\\ SHELS ($0<z<0.2$) & 412 & 11 & $2.304 \pm 0.007$ && $0.34 \pm 0.02$\\ SHELS ($0.2<z<0.3$) & 1115 & 11 & $2.277 \pm 0.005$ && $0.28 \pm 0.01$\\ SHELS ($0.3<z<0.4$) & 1316 & 11 & $2.277 \pm 0.004$ && $0.31\pm 0.01$\\ SHELS ($0.4<z<0.5$) & 818 & 11 & $2.299 \pm 0.005$ && $0.27\pm0.02$\\ SHELS ($0.5<z<0.6$) & 443 & 11 & $2.28 \pm 0.01$ && $0.26 \pm 0.03$\\ SHELS ($0.6<z<0.7$) & 61 & 11 & $2.24 \pm 0.07$ && $0.32 \pm 0.18$\\ \cutinhead{The volume limited sample (Figure \ref{fig:fit_rlim})}\\ SDSS VL & 109330 & 10.26 & $2.068 \pm 0.001$ & & $0.305 \pm 0.001$ \\ SHELS VL & 1827 & 10.26 & $2.06 \pm 0.01$ & & $0.30 \pm 0.01$ \\ SHELS VL ($0<z<0.2$) & 233 & 10.26 & $2.03 \pm 0.02$ && $0.37 \pm 0.03$\\ SHELS VL ($0.2<z<0.3$) & 641 & 10.26 & $2.05 \pm 0.01$ && $0.30 \pm 0.02$\\ SHELS VL ($0.3<z<0.4 $) & 953 & 10.26 & $2.03 \pm 0.01$ && $0.33\pm 0.02$\\ \cutinhead{The \citeauthor{Hyde2009a}, \citeauthor{Belli2014} and \citeauthor{Rines2016} samples (Figures \ref{fig:belli} and \ref{fig:rines})} \\ {Hyde+2009} & 46410 & 11 & $2.29\pm 0.02$ & & $0.286 \pm 0.002$ \\ Belli+2014 & 56 &11 & $2.38 \pm 0.01$ & & $0.30 \pm 0.03$ \\ Rines+2016 & 123 & 11 & $1.73 \pm 0.03$ & &$0.32 \pm 0.03$ \\ \enddata \tablecomments{Parameters of the fits. Column 1 identifies the sample and Column 2 gives the number of object in the sample. Column 3 is the ``break point" stellar mass. An error indicates that the break point is fit to the data; otherwise $M_b$ is fixed. Column 4 gives the velocity dispersion at $M_b$. Column 5 and 6 give the power law index of the fit. An entry for $\alpha_2$ alone indicates a single power law fit to the data.} \label{tab:fit_param} \end{deluxetable*}
16
7
1607.04275
We examine the relation between stellar mass and central stellar velocity dispersion—the M <SUB>*</SUB> σ relation—for massive quiescent galaxies at z &lt; 0.7. We measure the local relation from the Sloan Digital Sky Survey and the intermediate redshift relation from the Smithsonian Hectospec Lensing Survey. Both samples are highly complete (&gt;85%) and we consistently measure the stellar mass and velocity dispersion for the two samples. The M <SUB>*</SUB> σ relation and its scatter are independent of redshift with σ \propto {M}<SUB>* </SUB><SUP>0.3</SUP> for M <SUB>*</SUB> ≳ 10<SUP>10.3</SUP> M <SUB>⊙</SUB>. The measured slope of the M <SUB>*</SUB> σ relation is the same as the scaling between the total halo mass and the dark matter halo velocity dispersion obtained by N-body simulations. This consistency suggests that massive quiescent galaxies are virialized systems, where the central dark matter concentration is either a constant or negligible fraction of the stellar mass. The relation between the total galaxy mass (stellar + dark matter) and the central stellar velocity dispersion is consistent with the observed relation between the total mass of a galaxy cluster and the velocity dispersion of the cluster members. This result suggests that the central stellar velocity dispersion is directly proportional to the velocity dispersion of the dark matter halo. Thus, the central stellar velocity dispersion is a fundamental, directly observable property of galaxies, which may robustly connect galaxies to dark matter halos in N-body simulations. To interpret the results further in the context of ΛCDM, it would be useful to analyze the relationship between the velocity dispersion of stellar particles and the velocity dispersion characterizing their dark matter halos in high-resolution cosmological hydrodynamic simulations.
false
[ "dark matter halos", "stellar mass and central stellar velocity dispersion", "the dark matter halo velocity dispersion", "stellar particles", "the central stellar velocity dispersion", "massive quiescent galaxies", "galaxies", "the dark matter halo", "the stellar mass and velocity dispersion", "the velocity dispersion", "the central dark matter concentration", "redshift", "high-resolution cosmological hydrodynamic simulations", "N-body simulations", "the intermediate redshift relation", "the total galaxy mass", "the stellar mass", "their dark matter", "M}<SUB", "the Smithsonian Hectospec Lensing Survey" ]
11.72675
5.770826
-1
12405952
[ "Solarz, A.", "Takeuchi, T. T.", "Pollo, A." ]
2016A&A...592A.155S
[ "Total infrared luminosity estimation from local galaxies in AKARI all sky survey" ]
7
[ "National Center for Nuclear Research, ul. Hoża 69, 00-681, Warsaw, Poland", "Division of Particle and Astrophysical Science, Nagoya University, Furo-cho, Chikusa-ku, 464-8602, Nagoya, Japan", "National Center for Nuclear Research, ul. Hoża 69, 00-681, Warsaw, Poland; The Astronomical Observatory of the Jagiellonian University, ul. Orla 171, 30-244, Kraków, Poland" ]
[ "2017PASP..129d4102D", "2017arXiv170701652M", "2019ApJ...872...49F", "2019PASJ...71...28S", "2019csic.conf..203P", "2021ApJ...908..246Y", "2022arXiv220400831T" ]
[ "astronomy" ]
4
[ "infrared: galaxies", "galaxies: star formation", "galaxies: evolution", "Astrophysics - Astrophysics of Galaxies" ]
[ "1987ARA&A..25..187S", "1988ApJS...68..151H", "1992ifss.book.....M", "1996A&A...315L..27K", "1996ARA&A..34..749S", "1998ApJ...500..525S", "2000prpl.conf...97W", "2001ApJ...549..215D", "2002ApJ...576..159D", "2004ApJS..154....1W", "2004fost.book.....S", "2005A&A...440L..17T", "2005MNRAS.362..592T", "2005fds..book.....S", "2007ApJ...657..810D", "2007ApJS..173..404B", "2007PASJ...59S.389K", "2008PASJ...60S.477H", "2010A&A...514A...3P", "2010A&A...514A...4T" ]
[ "10.1051/0004-6361/201629062", "10.48550/arXiv.1607.08747" ]
1607
1607.08747_arXiv.txt
The question of how and when galaxies have formed their stellar content is paramount when addressing the topic of formation and evolution of galaxies. The baryonic gas contained within a dark matter halo goes through a cooling process, which results in a decrease of pressure. As a result, it starts to flow into the centre of the halos' potential well, creating an excess of density. When it exceeds the density of the dark matter, the gas starts to act under its own gravity and eventually collapses. This process ultimately leads to the creation of a dense cloud of molecular gas, which will serve as a 'nursery' for the star formation (e.g. \citealt{williams}). The process of formation of stars in galaxies is always accompanied by the dust production (e.g. \citealt{schulz}, \citealt{stahler}); however, the exact physics of this process is still a subject of investigation (e.g. \citealt{ttt05}). One thing is certain, though: the dust particles in galaxies absorb the ultra-violet (UV) radiation emitted by young, hot stars, and re-emit it in the far infrared (FIR). Therefore, the observations at these wavelengths are in tight correlation with the amount of the dust contained in a galaxy and with its stellar content properties. Majority of the star forming regions within a galaxy are hidden in dense gas clouds (e.g. \citealt{sasiadka}), which makes them impossible to observe in the UV. This fact is also confirmed by analyses of the evolution of the luminosity function (the dependence of the number of stars/galaxies on their absolute magnitude) both in FIR and UV \citep{ttt05a}, with the former one being far stronger than the latter, which means that the amount of the hidden star formation zones grows proportionally with redshift (for $z<1$). All those reasons make the infrared observations of the Universe crucial for understanding the star formation history of the Universe. One of the ways to quantify star-formation is to measure the total infrared luminosity ($L_{\textrm{TIR}}$), as the two are tightly correlated (e.g. \citealt{helou88}, \citealt{sanders96}, \citealt{dale01}, \citealt{dale02}, \citealt{draine07}, \citealt{takeuchi10}). Therefore, by using the emission measured throughout different IR passbands it is possible to recover the information about the star-formation activity of galaxies. The purpose behind the launch of AKARI satellite was to make all-sky surveys at infrared wavelengths with better sensitivity, spatial resolution and wavelength coverage than that of its predecessor: pioneering IRAS \citep{soi}. Other past IR satellites, ISO \citep{kess} and Spitzer Space Telescope \citep{wer}, though having better capabilities than IRAS, were not designed for all-sky surveys. The AKARI satellite was launched by JAXA’s MV8 vehicle on February 22, 2006, and one of its main goals was to perform and all-sky survey at far-infrared (FIR) wavelengths. For this purpose a dedicated FIS camera \citet{kawada07} was used. It covered the wavelength range between 50 and 180~$\mathrm{\mu}$m through exposure of 4 FIR filters: \textit{N60} centred at 65~$\mathrm{\mu}$m, WIDE-S centred at 90~$\mathrm{\mu}$m, WIDE-L centred at 140~$\mathrm{\mu}$m and \textit{N160} centred at 160~$\mathrm{\mu}$m. In this work, we use new and secure data from AKARI all sky survey to present a recalibrated formula to recover $L_{\textrm{TIR}}$. We adopt a cosmological model with $\Omega_{m}$ = 0.3, and $\Omega_{\Lambda}$ = 0.7, $H_{0}=70$~km~$\mathrm{s^{-1}}$~$\mathrm{Mpc^{-1}}$ and define $L_{\nu}$~$[\mathrm{erg}$~$\mathrm{s^{-1}}$~$\mathrm{Hz^{-1}}]$ as the luminosity per unit frequency at a frequency $\nu = c/\lambda$ (c: the speed of light) throughout this paper.
In this work we used the new release of the AKARI all sky survey catalogue performed at far-infrared and covering the wavelength range from 50 to 180 $\mathrm{\mu}$m to recalibrate the formula to obtain total infrared luminosities. We find very tight correlations between the luminosity estimates from AKARI FIS catalogue and luminosity estimates from IRAS FSC catalogue. We do not find any significant enhancement of the correlations when only two wide filters are used instead of three AKARI-FIS filter combination ($N60$, $\mbox{WIDE-S}$, $\mbox{WIDE-L}$).
16
7
1607.08747
We aim to use the a new and improved version of AKARI all sky survey catalogue of far-infrared sources to recalibrate the formula to derive the total infrared luminosity. We cross-match the faint source catalogue (FSC) of IRAS with the new AKARI-FIS and obtained a sample of 2430 objects. Then we calculate the total infrared (TIR) luminosity L<SUB>TIR</SUB> from the formula previously estimated from IRAS data and compare it with total infrared luminosity from AKARI FIS bands to obtain new coefficients for the general relation to convert FIR luminosity from AKARI bands to the TIR luminosity.
false
[ "total infrared luminosity", "AKARI FIS bands", "FIR luminosity", "AKARI bands", "AKARI FIS", "AKARI", "new coefficients", "IRAS data", "the total infrared luminosity", "IRAS", "TIR", "the TIR luminosity", "TIR</SUB", "AKARI-FIS", "FIR", "the new AKARI-FIS", "all sky survey catalogue", "the faint source catalogue", "FSC", "the general relation" ]
13.413519
7.744332
132
2555340
[ "González-López, J.", "Bauer, F. E.", "Romero-Cañizales, C.", "Kneissl, R.", "Villard, E.", "Carvajal, R.", "Kim, S.", "Laporte, N.", "Anguita, T.", "Aravena, M.", "Bouwens, R. J.", "Bradley, L.", "Carrasco, M.", "Demarco, R.", "Ford, H.", "Ibar, E.", "Infante, L.", "Messias, H.", "Muñoz Arancibia, A. M.", "Nagar, N.", "Padilla, N.", "Treister, E.", "Troncoso, P.", "Zitrin, A." ]
2017A&A...597A..41G
[ "The ALMA Frontier Fields Survey. I. 1.1 mm continuum detections in Abell 2744, MACS J0416.1-2403 and MACS J1149.5+2223" ]
63
[ "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile; Millennium Institute of Astrophysics, 1515 Las Condes, Santiago, Chile; Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO, 80301, USA", "Millennium Institute of Astrophysics, 1515 Las Condes, Santiago, Chile; Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile; Núcleo de Astronomía, Facultad de Ingeniería, Universidad Diego Portales, 441 Av. Ejército, Santiago, Chile", "Joint ALMA Observatory, 3107 Alonso de Córdova, Vitacura, Santiago, Chile; European Southern Observatory, 3107 Alonso de Córdova, Vitacura, 19001 Casilla, Santiago, Chile", "Joint ALMA Observatory, 3107 Alonso de Córdova, Vitacura, Santiago, Chile; European Southern Observatory, 3107 Alonso de Córdova, Vitacura, 19001 Casilla, Santiago, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile; Department of Physics and Astronomy, University College London, Gower Street, London, WC1E 6BT, UK", "Departamento de Ciencias Físicas, Universidad Andres Bello, 252 Avenida República, Santiago, Chile; Millennium Institute of Astrophysics, 1515 Las Condes, Santiago, Chile", "Núcleo de Astronomía, Facultad de Ingeniería, Universidad Diego Portales, 441 Av. Ejército, Santiago, Chile", "Leiden Observatory, Leiden University, 2300 RA, Leiden, The Netherlands", "Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD, 21218, USA", "Universität Heidelberg, Zentrum für Astronomie, Institut für Theoretische Astrophysik, Philosophenweg 12, 69120, Heidelberg, Germany", "Department of Astronomy, Universidad de Concepcion, Casilla 160-C, Concepción, Chile", "Department of Physics and Astronomy, Johns Hopkins University, Baltimore, MD, 21218, USA", "Instituto de Física y Astronomía, Universidad de Valparaíso, Avda. Gran Bretaña 1111, Valparaíso, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Instituto de Astrofísica e Ciências do Espaço, Universidade de Lisboa, OAL, Tapada da Ajuda, 1349-018, Lisboa, Portugal", "Instituto de Física y Astronomía, Universidad de Valparaíso, Avda. Gran Bretaña 1111, Valparaíso, Chile; Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Department of Astronomy, Universidad de Concepcion, Casilla 160-C, Concepción, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Instituto de Astrofísica and Centro de Astroingeniería, Facultad de Física, Pontificia Universidad Católica de Chile, Casilla 306, 22, Santiago, Chile", "Cahill Center for Astronomy and Astrophysics, California Institute of Technology, MC 249-17, Pasadena, CA, 91125, USA; Physics Department, Ben-Gurion University of the Negev, PO Box 653, 84105, Be'er-Sheva, Israel" ]
[ "2016ApJ...833...25Z", "2017A&A...602A..54M", "2017A&A...604A.132L", "2017A&A...608A.138G", "2017ApJ...837..139C", "2017ApJ...840...29H", "2017ApJ...846L..22G", "2017ApJ...850...83F", "2017ApJ...850..189H", "2017ConPh..58..160B", "2018A&A...620A.125M", "2018A&A...620A.152F", "2018ApJ...854...83T", "2018ApJ...861....7F", "2018ApJ...868..115Y", "2018ApJS..235...14S", "2018Natur.557..392H", "2018PASJ...70..105H", "2019A&A...625A..96K", "2019ApJ...871..109P", "2019ApJ...874...27T", "2019ApJ...882..136A", "2019ApJ...882..139G", "2019MNRAS.490.4956G", "2019PASJ...71...71H", "2020A&A...633A.160C", "2020A&A...643A...2B", "2020ApJ...888...44C", "2020ApJ...897...91G", "2020ApJ...901...79A", "2020IAUGA..30..423M", "2020MNRAS.493.4294B", "2020PASJ...72...69Y", "2020RSOS....700556H", "2020arXiv200810641Z", "2021ApJ...908..192S", "2021ApJ...910..106J", "2021ApJ...919...54Z", "2022A&A...658A..43G", "2022A&A...658A..77N", "2022ApJ...927..141M", "2022ApJ...932...77S", "2022ApJ...939....5C", "2022ApJS..263...38K", "2022arXiv221204026B", "2023A&A...675A..85M", "2023ApJ...942L...1W", "2023ApJ...945..121U", "2023ApJ...952...28C", "2023ApJ...952..142F", "2023ApJ...956...95B", "2023ApJ...957...63B", "2023MNRAS.523.4568F", "2023MNRAS.526.2423G", "2023arXiv230301658F", "2024A&A...684A.119P", "2024A&A...685A.161R", "2024ApJ...962..128M", "2024ApJ...965..108U", "2024Natur.628...57F", "2024SSRv..220...19N", "2024arXiv240205849U", "2024arXiv240609890T" ]
[ "astronomy" ]
12
[ "gravitational lensing: strong", "galaxies: high-redshift", "submillimeter: galaxies", "Astrophysics - Astrophysics of Galaxies" ]
[ "1996MNRAS.283.1340B", "1997ApJ...490L...5S", "1999ASPC..180..127B", "2000AJ....119.2092B", "2001ITAP...49.1683P", "2002AJ....123.2197C", "2002PhR...369..111B", "2004ApJ...616...71S", "2004ApJ...616L...1H", "2004ApJS..154..193W", "2005A&A...437...39B", "2005ApJ...622..772C", "2005ApJ...632..736A", "2006MNRAS.367.1209L", "2007ASPC..376..127M", "2007ApJ...660L..43N", "2007MNRAS.377.1557N", "2008ApJ...688...59Y", "2009A&A...500..681M", "2009ApJ...706.1201B", "2009MNRAS.396.1985Z", "2010A&A...518L..44K", "2010MNRAS.407.1268Y", "2010Natur.464..733S", "2010Sci...330..800N", "2011A&A...528A..35M", "2011MNRAS.415.3081I", "2011MNRAS.415.3473V", "2011MNRAS.417..333M", "2012ApJ...761...20H", "2012MNRAS.424.2429S", "2013ApJ...762...59W", "2013ApJ...762L..30Z", "2013ApJ...779...25B", "2013ApJ...779...32V", "2013PASP..125..306F", "2014ApJ...785..111W", "2014ApJ...786...60A", "2014ApJ...788..125S", "2014ApJ...793L..12Z", "2014ApJ...795...93Z", "2014ApJ...797...48J", "2014MNRAS.439.2096W", "2014MNRAS.444..268R", "2014PhR...541...45C", "2015A&A...575A..92L", "2015A&A...577A..29M", "2015ApJ...798L..18H", "2015ApJ...799...81S", "2015ApJ...800...84C", "2015ApJ...808L...4A", "2015ApJ...812...43B", "2015ApJ...812..114T", "2015ApJ...815...18I", "2015MNRAS.446.4132J", "2015MNRAS.447.3130D", "2015MNRAS.450.3032M", "2015MNRAS.451L..40R", "2015MNRAS.452.1437J", "2016ApJ...819..114K", "2016ApJ...826..112S", "2016ApJS..222....1F", "2016MNRAS.457.2029J", "2016MNRAS.459.1626R", "2016MNRAS.461.2126S", "2017ApJ...837...97L", "2017MNRAS.465.1030P", "2017MNRAS.466..861D" ]
[ "10.1051/0004-6361/201628806", "10.48550/arXiv.1607.03808" ]
1607
1607.03808_arXiv.txt
\hspace{1cm}Past studies of dusty, star-forming galaxies (DSFGs) at infrared (IR) through radio wavelengths have firmly established their role in the growth and evolution of massive galaxies across cosmic time \citetext{see reviews by \citealp{Blain2002} and \citealp{Casey2014}}. At the bright end, DSFGs are observed to have IR luminosities in excess of $10^{13}$ $L_{\odot}$, yet appear extraordinarily compact, with typical resolved half-light radii of $R_{\rm e}$$\approx$1--1.5\,kpc \citetext{e.g., \citealp{Bussmann2013,Ikarashi2014,Simpson2015,Miettinen2015}; hereafter B13, I14, S15 and M15, respectively}. The bulk of their IR emission is thought to be powered by star formation \citep{Alexander2005}, with unobscured star-formation rates (SFRs) of up to several thousands of $M_{\odot}$\,yr$^{-1}$ and SFR densities of $\gtrsim$100 $M_{\odot}$\,yr$^{-1}$\,kpc$^{-2}$. As such, DSFGs represent the most intense starbursts in the Universe. While little is known about fainter DSFGs (sub-mJy population), extrapolation of the DSFG population is estimated to account for roughly half of the IR background light in aggregate \citep[e.g.,][]{Magnelli2011, Viero2013}. Obtaining a working knowledge of the underlying physics that drives the distribution of faint lensed DSFGs thus appears to be critical for understanding cosmic galaxy assembly. Many DSFGs are optically faint ($I_{AB}$$\gtrsim$24\,mag, $K_{AB}$$\sim$21-22), due to a combination of high redshift and strong dust-extinction \citep[e.g.,][]{Barger2000, Smail2004, Chapman2005}. This has made unbiased estimates of their population statistics and physical properties (i.e., dust, gas and stellar contents and morphologies) challenging and expensive with current technology. It also suggests that finding and characterizing the fainter DSFG population may prove even harder. Nonetheless, the fainter population is particularly interesting because their SFRs are in the same range as those found for ultraviolet (UV) and optically selected samples such as star-forming $BzK$ galaxies, BX/BM galaxies, and Lyman break galaxies (LBGs), which comprise the normal galaxy ‘main sequence’ \citep[e.g.,][]{Noeske2007}. Meaningful comparisons between the more abundant faint DSFGs and these unobscured main sequence populations could help to elucidate the factors that determine the dust content in galaxies with comparable properties. One way to make progress toward the study of faint DSFGs, in the face of instrumental confusion limits at far-IR (FIR) and submillimeter (submm) wavelengths ($\sim$2--7\,mJy beam$^{-1}$ between 0.25--1.3\,mm) as well as traditional difficulties associated with expensive multi-wavelength follow-up, is to leverage with the power of gravitational lensing. Such studies have helped provide detailed characterizations of many dozens of intrinsically faint DSFGs selected either in wide area submm surveys as galaxy-galaxy lenses \citep[e.g.,][]{Blain1996, Negrello2007, Negrello2010, Wardlow2013} or behind massive lensing galaxy clusters \citep[e.g.,][]{Smail1997, Cowie2002, Swinbank2010}. Although immensely insightful, such studies can suffer from additional uncertainties due to the quality of the mass models, microlensing, and potential biases stemming from the cross-section of the DSFG population which is more easily lensed (i.e., compact starbursts). Robustly determined mass models and observations over long timescales can limit the impact of these problems. Building on past works of lensed DSFGs, here we employ the novel sensitivity of the Atacama Large Millimeter/submillimeter Array (ALMA) to detect DSFGs roughly 1 dex fainter than the aforementioned confusion limits, behind three strong-lensing galaxy clusters: Abell 2744 ($z=0.308$), MACSJ0416.1-2403 ($z=0.396$) and MACSJ1149.5+2223 ($z=0.543$) (hereafter A2744, MACSJ0416, MACSJ1149, respectively).\footnote{Approved ALMA Cycle 3 observations of the remaining three FFs clusters, namely MACSJ0717.5+3745, Abell 370, and Abell 1063S, are in progress.} Importantly, these clusters are part of the Frontier Fields (FFs) Survey,\footnote{\url{http://www.stsci.edu/hst/campaigns/frontier-fields/}} a legacy project which combines the power of gravitational lensing by massive clusters (with magnifications of $\mu$$>$5--10 over up to several arcmin$^{2}$ regions and 100's of multiple images) with extremely deep multi-band {\it HST} and {\it Spitzer} imaging of six lensing clusters and adjacent parallel fields \citep{Coe2015}. While the primary goal of these observations is to potentially detect and characterize $z\geq1$ galaxies 10--50 times intrinsically fainter than any seen before \citep[e.g.,][]{Atek2014, Zitrin2014, Zheng2014, Laporte2015, Mcleod2015, Infante2015}, the {\it HST}+{\it Spitzer} observations and their associated ancillary data enable several ALMA related science cases. The FFs campaign is comprised of allocations of 840 {\it HST} orbits and 1000 {\it Spitzer} hrs, respectively, taking advantage of the two Great Observatories unsurpassed spatial resolution and/or depth. The fields have amassed a wealth of multi-wavelength ancillary data. Notably, each cluster already has extensive space+ground-based archival data (e.g., 16-band {\it HST}, {\it XMM-Newton}, {\it Chandra}, {\it Spitzer}, {\it Herschel}, and VLA imaging) and more than 800 spectroscopically confirmed sources (cluster+background galaxies), including VLT/VIMOS confirmations for most of the major gravitational arcs/images down to $\sim$26 ABmag (PI Rosati). To this, the FFs campaign adds deep ACS ($F435W$$=$$F606W$$=$$F814W$$\approx$28.4--29.0\,ABmag), WFC3 ($F105W$$=$$F125W$$=$$F140W$$=$$F160$$\approx$29.1--29.4\,ABmag), and IRAC1/IRAC2 ($\approx$25.0\,ABmag) imaging \citep[e.g.,][]{Coe2015, Lotz2016}. These are complemented by new and growing data from {\it HST}, {\it Chandra}, JVLA, VLT HAWK-I/VIMOS/MUSE, Keck DEIMOS/LRIS/MOSFIRE, and Magellan FourSTAR/IMACS/MMIRS. The current data already allow the assembly of 10,000's of accurate $z_{\rm ph}$'s, 400/200/10's of Lyman Break dropouts at $z$$>$6/7/8, and spectacularly resolved images of $>$600 $z$$=$1--6 gravitational arcs (2--10") and multiple-lenses in a central region of each cluster \citep[e.g.,][]{Richard2014}. The FFs clusters represent the best studied gravitational lenses, with the best available estimates of magnifications and uncertainties (ever-improving lens models are available from many teams).\footnote{\url{http://www.stsci.edu/hst/campaigns/frontier-fields/Lensing-Models}} Thus they are key regions on the sky to observe normal galaxies at high-$z$, and warrant strong mm constraints. The search for DSFGs on the FFs using {\it Herschel} data has already found nine sources with magnification $\mu\geq4$ at $z<1.5$ and other nine sources at $z>2$ \citep{Rawle2016}. Each of the three FFs clusters is being complemented with a $\approx$2\farcm1$\times$2\farcm2 1.1\,mm ALMA mosaic, coincident with the deep {\it HST} WFC3/IR imaging region, achieving an rms depth of 55--71 $\mu$Jy\,${\rm beam}^{-1}$, depending on the field. The 1.1\,mm band benefits from the negative K-correction due to the typical galaxy spectral energy distribution (SED), and thus effectively provides a redshift unbiased census between $z$$\sim$1--10 \citep[e.g.,][]{Blain2002}. The data allow us to pinpoint locations of extreme star formation and probe factors of 2--10$\times$ fainter into the $L_{\rm IR}$ and SFR distributions compared to similar blank-field surveys due to typical lensing boosts. The data also allow constraints on the SFRs and line emission for 1000's of optical and IR objects undetected in our 1.1\,mm survey (individually and in aggregate). This paper is the first in a series, and primarily describes the reduction and analysis of the ALMA 1.1\,mm data for A2744, MACSJ0416, and MACSJ1149 (hereafter, the ALMA-FFs). We focus this paper on the observed sources properties in the lensed image plane. To fully characterize the detected sources in the source plane, a combination of redshifts estimates and mass models for the galaxy clusters is needed. Separate companion papers will detail the identification, redshift estimation and initial characterization of multi-wavelength counterparts (N. Laporte et al. 2016, in prep.), physical properties based on SED-fitting and morphological studies in the source plane (J. Gonz\'alez-L\'opez et al. 2016, in prep.), number counts (A. Mu{\~{n}}oz Arancibia et al. 2016, in prep.), and stacking analyses of multiply imaged sources and dropout candidates (R. Carvajal et al. 2016, in prep.). This paper is organized as follows: in $\S$2 we outline the observations and describe the reduction and imaging procedures; in $\S$3 we discuss our source detection methods and sensitivity estimates; in $\S$4 we assess the fluxes and angular extents of the detected sources; and in $\S$5 we summarize our results. Throughout the paper we assume a concordance cosmology and quote errors at 1$\sigma$ unless stated otherwise.
\label{sec:discuss} \begin{table*} \caption[]{Flux density measurements for high-significance continuum detections. {\it Col. 1}: Source ID. {\it Col. 2}: Integrated flux density and 1$\sigma$ statistical error assuming a PS model, in mJy. {\it Col. 3}: Integrated flux density and 1$\sigma$ statistical error assuming a 6-parameter extended Gaussian model (EXT), in mJy. {\it Col. 4}: Reduced $\chi^{2}/DOF$ for PS model. {\it Col. 5}: Reduced $\chi^{2}/DOF$ for EXT model. {\it Col. 6}: Integrated flux density and 1$\sigma$ statistical error from $uv$ fitting, in mJy. \label{tab:gold_flux}} \centering \begin{tabular}{cccccc} \hline {ID}&{F$_{\rm Int,\,PS}$ [mJy]} & {F$_{\rm Int,\,EXT}$ [mJy]} & {$\chi^2_{\rm red,\,PS}$} & {$\chi^2_{\rm red,\,EXT}$} & {F$_{\rm uv-fit}$ [mJy]} \\ \hline \noalign{\smallskip} A2744-ID01 &$ 1.495 \pm 0.081 $&$ 1.545 \pm 0.081 $& 1.7 & 1.6 &$1.570\pm0.073$\\ A2744-ID02 &$ 1.656 \pm 0.131 $&$ 3.262 \pm 0.213 $& 10.3 & 2.2 &$2.816\pm0.229$\\ A2744-ID03 &$ 1.074 \pm 0.084 $&$ 1.979 \pm 0.137 $& 7.9 & 2.0 &$1.589\pm0.125$\\ A2744-ID04 &$ 1.012 \pm 0.099 $&$ 1.173 \pm 0.111 $& 2.0 & 1.7 &$1.009\pm0.074$\\ A2744-ID05 &$ 0.859 \pm 0.125 $&$ 1.422 \pm 0.188 $& 3.0 & 0.7 &$1.113\pm0.135$\\ A2744-ID06 &$ 0.822 \pm 0.129 $&$ 2.274 \pm 0.274 $& 12.1 & 11.9 &$1.283\pm0.241$\\ A2744-ID07 &$ 0.500 \pm 0.108 $&$ 0.716 \pm 0.149 $& 1.5 & 1.4 &$0.539\pm0.082$\\ MACSJ0416-ID01 &$ 1.144 \pm 0.089 $&$ 1.356 \pm 0.111 $& 3.0 & 0.6 &$1.319\pm0.103$\\ MACSJ0416-ID02 &$ 0.455 \pm 0.081 $&$ 0.626 \pm 0.111 $& 2.8 & 2.0 &$0.574\pm0.132$\\ MACSJ0416-ID03 &$ 0.375 \pm 0.091 $&$ 0.477 \pm 0.114 $& 4.7 & 4.3 &$0.411\pm0.072$\\ MACSJ0416-ID04 &$ 0.367 \pm 0.089 $&$ 0.456 \pm 0.111 $& 3.0 & 2.9 &$0.478\pm0.166$\\ MACSJ1149-ID01 &$ 0.500 \pm 0.104 $&$ 0.635 \pm 0.125 $& 1.1 & 0.5 &$0.579\pm0.134$\\ \hline \end{tabular} \end{table*} \begin{figure*}[!htbp] \includegraphics[width=\hsize/2]{plots/A2744-ID01.pdf} \includegraphics[width=\hsize/2]{plots/A2744-ID02.pdf} \includegraphics[width=\hsize/2]{plots/A2744-ID03.pdf} \includegraphics[width=\hsize/2]{plots/A2744-ID04.pdf} \includegraphics[width=\hsize/2]{plots/A2744-ID05.pdf} \includegraphics[width=\hsize/2]{plots/A2744-ID06.pdf} \includegraphics[width=\hsize/2]{plots/A2744-ID07.pdf} \includegraphics[width=\hsize/2]{plots/MACSJ0416-ID01.pdf} \includegraphics[width=\hsize/2]{plots/MACSJ0416-ID02.pdf} \includegraphics[width=\hsize/2]{plots/MACSJ0416-ID03.pdf} \caption{Cutout images of the secure continuum detections in the FFs clusters. \textit{left:} 1.1\,mm continuum emission image adopting natural weighting, with colors as in Figure~\ref{fig:map_A2744}. Solid black curves show the positive $S/N$ contours with natural weighting, starting from $\pm2\sigma$ up to $\pm5\sigma$ in steps of $1\sigma$, and then above $\pm5\sigma$ in steps of $2.5\sigma$. Dashed curves show the negative $S/N$ contours. In the bottom left corner we show the corresponding synthesized beam. For A2744-ID06 we additionally overlay solid purple curves showing the $S/N$ contours from the corresponding Taper weighted image. The best-fit 2-d elliptical Gaussian model is shown as a green region whose size denotes where the emission is half of the maximum. The yellow points represent the position of optical/NIR detected galaxies. \textit{right:} Residual of the 1.1\,mm continuum emission after the best-fit model is subtracted from the $uv$ visibilities. The color scale is identical to the left-hand side. \label{fig:cont_fits1}} \end{figure*} \begin{figure*}[!htbp] \includegraphics[width=\hsize/2]{plots/MACSJ0416-ID04.pdf} \includegraphics[width=\hsize/2]{plots/MACSJ1149-ID01.pdf} \caption{Continuation of Figure~\ref{fig:cont_fits1}. \label{fig:cont_fits2}} \end{figure*} \begin{figure*}[!htbp] \includegraphics[width=\hsize/3]{plots/A2744-ID01_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/A2744-ID02_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/A2744-ID03_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/A2744-ID04_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/A2744-ID05_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/A2744-ID06_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/A2744-ID07_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/MACSJ0416-ID01_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/MACSJ0416-ID02_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/MACSJ0416-ID03_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/MACSJ0416-ID04_counterpart.pdf} \includegraphics[width=\hsize/3]{plots/MACSJ1149-ID01_counterpart.pdf} \caption{6\arcsec$\times$6\arcsec color image cutouts of the NIR counterparts to the ALMA detections, with blue corresponding to bands $F435W$$+$$F060W$, green to bands $F810W$$+$$F105W$ and red to bands $F125W$$+$$F140W$$+$$F160W$. The cyan contours denote the ALMA 1.1\,mm emission, whereby the levels are displayed on a logarithmic scale at $\pm$ 3, 4, 5.3, 7.1, 9.5, 12.6, 16.8, 22.5, 30 $\times\sigma$, with $\sigma$ being the primary beam corrected noise level at the position of the source. The ALMA synthesized beam corresponding to each observation is shown in the lower left corner of each cutout. A green star represents the position of the associated NIR counterpart, while the magenta ellipse shows the $1\sigma$ range of the center of the ALMA emission based on the $uv$-plane fit. In all cases, the positional error associated with the NIR position is small compared to that of the ALMA emission. \label{fig:counterparts}} \end{figure*} \subsection{Purity}\label{sec:purity} In order to establish the purity levels of our detections, we need to estimate how many false detections we expect per $S/N$ for a given observation. We use the task \texttt{Simobserve} in {\sc casa} to simulate observations with similar properties as the ones obtained in our campaign. For the simulations the thermal noise is added using an atmospheric profile for the ALMA site \citep{Pardo2001} and assuming the observations are performed when the field transits (this assumption is not critical for the purposes here). We adopt the nominal second (MACSJ0416 and MACSJ1149) and third (A2744) most compact ALMA antenna configurations (C36-2 and C36-3) during cycle 2, which are the closest matches to the ones used in the real observations. The other parameters such as frequency coverage, PWV and integration time are set to achieve an rms similar to the ones obtained for each cluster field. After construction of the simulations, we used the noisy measurement set files to create images in the same manner as used to produce the real observations. These images were then analyzed with the same algorithm used to detect sources (\S\ref{sec:imaging}). We performed 30 simulations per antenna configuration and recorded the number of detections with $S/N\geq5$ recovered by our code. The number of detections was then renormalized as a function of $S/N$ for each antenna configuration. For all antenna configurations, we obtain $N\ll1$ simulated-noise detections for $S/N\geq5$, meaning that our $S/N\geq5$ sample should all be true detections assuming that the noise properties of the simulations are similar to those in the observations. The deviation from a perfect Gaussian distribution ratio at $S/N$$\sim$$-5.0$ in the data distribution observed in Figure \ref{fig:sn_clusters} and discussed in $\S$\ref{Sec:high_significance_cont_detections} can be explained by low-number statistics, as the same behavior is observed in the simulated observations histograms when Gaussian noise is assumed. It is important to note that because of the stochastic behavior in the extreme tails of the distribution, extra care needs to be taken when using the negative pixels as a noise reference, as some fraction of noise distributions will never be completely symmetrical, even when Gaussian noise is assumed. This effect can easily lead to misinterpretations when using the negative pixels as a reference to estimate the purity of continuum detections. For instance, in MACSJ0416 and MACSJ1149, we find no negative sources with $S/N$$\leq$$-5.0$, while in A2744 we find two sources with $S/N$$\leq$$-5.0$ compared to seven sources with $S/N$$\leq$$+5.0$. The implied purity for A2744 is thus $\sim$0.7 based on negative count symmetry. However, based on the NIR counterparts found for all the ALMA sources with $S/N\geq5$, we estimate a "true" detection purity closer to $\approx$1 for the three clusters, consistent with the simulations. Although the difference between the purity estimated from negative counts and the "true" purity is not statistically significant (only $1.1\sigma$), it exemplifies the extra care needed when using the negative counts for purity estimates. \subsection{Flux estimates}\label{sec:flux} \subsubsection{Image fit} We obtained three flux density measurements for each source. The first two methods fit for the flux density in the image plane assuming a point source (PS) and an extended source (6 parameters Gaussian function, hereafter EXT), respectively, while the third method fits for the flux density in the $uv$-plane. Fitting the sources in the image plane is substantially more straightforward, but inherits certain dependencies based on the weighting method used to construct the images. To measure the flux densities from the images, we fit an elliptical Gaussian to each source. For a PS, we fixed the size and angle of the elliptical Gaussian to that of the synthesized beam, and only allow for changes in flux and position. For an EXT, we allow all the six parameters describing the 2-d Gaussian to vary. For both fits, a reduced $\chi^2$ is estimated using all available image pixels that satisfy $S/N>2$ (to limit the fit to pixels associated with the source). The degrees of freedom are estimated as $DOF = N-P$, where $N$ is the number of pixels with $S/N>2$ and $P$ is the number of free parameters in the fit (three for a PS and six for an EXT). The measured flux densities and reduced $\chi^2$ values are presented in Table~\ref{tab:gold_flux}. These values were measured in the natural weighting images. The $1\sigma$ uncertainties in the flux densities measurements are of the order of 0.1 mJy. We should expect a number density of $\approx10^{4.8}\,\,{\rm deg}^{-2}$ sources with flux densities $>0.084$ mJy \citep{Fujimoto2016}. Even for the largest beam obtained in MACSJ1149 (Table~\ref{tab:imaging_results}) we obtain $\sim140$ beams per source at such flux density levels, therefore expecting a negligible flux boosting from confused sources in our results. \subsubsection{$uv$-plane fit} To measure the flux in the $uv$-plane we used \texttt{uvmcmcfit},\footnote{\url{https://github.com/sbussmann/uvmcmcfit}} which is a \texttt{Python} implementation to fit models to interferometric data in the $uv$ plane. The code allows one to extract the maximum amount of information from the observations, particularly if the sources appear marginally resolved. \texttt{uvmcmcfit} is designed to allow for de-lensing of sources in the case of galaxy-galaxy or galaxy-group magnification. We chose to first model the observed sources as if no magnification were present. This approach should return the shape and flux of the sources in the lensed image plane. Reconstructions to the source plane of each detected galaxy are beyond the scope of the current work and will be presented in J. Gonzalez-Lopez et al. (2016, in prep.). For each source, we adopt the simplest model, assuming that the FIR emission can be fit with an elliptical Gaussian described by the following parameters: the total intrinsic flux density, the position of the source (RA,DEC), the effective radius defined as $r_{s}=\sqrt{a_{s}\times b_{s}}$ ($b$ and $a$ are the major and minor axis respectively), the axial ratio ($q_{s}=b_{s}/a_{s}$), and the position angle in degrees east of north ($\phi_{s}$). The effective radius is a good measurement of the total size of the sources. The goodness of fit for a given set of model parameters is determined from the maximum likelihood estimate $L$ given by \begin{equation} L = \sum_{u,v}\left(\frac{|V_{\rm ALMA}-V_{\rm model}|^2}{\sigma^2}+\log{2\pi\sigma^2}\right), \end{equation} \noindent where $\sigma$ is the $1\sigma$ uncertainty level for each $uv$ complex visibility (hereafter visibility; Fourier transform of the intensity distribution on the sky) given by the associated weights. In practice, each source is typically observed by only a handful of pointings from the full mosaic, and each of these pointings will have an associated weight that we must apply to its visibilities to combine with those from neighboring pointings. Because well-determined weights for the visibilities of each pointing are needed, we performed the following procedure for each source before fitting. We estimate an initial source position based on the continuum images and select all pointings which lie within 19\arcsec (PB$\geq$0.1) of the initial position. We then shift the phase center of all the selected pointings to the position of the source. We correct the amplitudes of the visibilities by the initial PB correction of each pointing with respect to the position of the source as ${\rm data}_{i} = {\rm data}_{i}/{\rm PB}_{i}$, where $i$ is the index of each pointing. The weights of the visibilities are then corrected by the same PB correction as for the amplitudes as ${\rm weight}_{i} = {\rm weight}_{i}\times{\rm PB}^{2}_{i}$. We concatenate all the pointings into a new dataset, which should now contain a correct set of relative weights for the visibilities, since the appropriate factor was applied to the weights given by the calibration. Finally, we scale the weights such that: \begin{equation} \sum {\rm weight}\times{\rm real}^{2} + {\rm weight}\times{\rm Im}^2 = N_{\rm visibilities}. \end{equation} \texttt{uvmcmcfit} uses \texttt{EMCEE} \citep{Foreman-mackey2013} to sample the posterior probability density function (PDF) of the model parameters. Between 10,000--240,000 iterations were required, depending on the speed of convergence. We used a 'burn-in' phase of 5,000 iterations to identify the 'best-fit' model parameters. The flux density given by the posterior PDF is presented in the last column of Table~\ref{tab:gold_flux}. The errors presented correspond to the 1 $\sigma$ range of the posterior PDF. The best-fit models are presented in Figures~\ref{fig:cont_fits1} and \ref{fig:cont_fits2}. To test that we are recovering most of the observed flux density with the $uv$-plane fit, we created a "Taper" image for each observation. This Taper image also adopts natural weighting but includes a uvtaper with an outertaper equal to 1\farcs5, which yields a beam size $\gtrsim$1\farcs5 but substantially worse rms sensitivity. We apply the same method used to measure the flux density in the natural weighted images on the Taper images. In Fig. \ref{fig:fluxes_comparison} we present the flux density measured using the $uv$-plane fit in the natural image and the F$_{\rm Int,\,EXT}$ measured in the Taper images. We see that the two estimates agree within the errors, demonstrating that we apparently recover a substantial amount (if not all) of the extended flux with the $uv$-plane fit method. We recover the total flux density even in the complex case of A2744-ID06 (see Fig. \ref{fig:cont_fits1}), where the Taper image emission is substantially more extended than the natural weighted emission. Given these findings, we adopt the F$_{uv-fit}$ flux estimates for all sources. \begin{figure}[!htbp] \centering \includegraphics[width=\hsize]{plots/flux_comparison_taper.pdf} \caption{Comparison of the flux density measured using the $uv$-plane fit method with that measured in Taper images. We see that the $uv$-plane fit recovers most of the emission detected in the lower resolutions images, within the corresponding errors. \label{fig:fluxes_comparison}} \end{figure} \subsection{Source sizes}\label{sec:size} \begin{table*} \caption[]{Source size measurements for high-significance continuum detections. {\it Col. 1}: Source ID. {\it Col. 2}: Effective radius defined as $r_{\rm s}=\sqrt{a_{\rm s}\times b_{\rm s}}$, in arcseconds. {\it Col. 3}: The axial ratio, $q_{\rm s}=b_{\rm s}/a_{\rm s}$. {\it Col. 4}: The position angle in degrees east of north, $\phi_{\rm s}$. {\it Col. 5}: Magnification value ($\mu$) estimated using the available lensing models and assuming $z=2\pm1$ for the sources without spectroscopic redshift. For MACSJ0146-ID01 and MACSJ0146-ID02, we adopt $z_{\rm spec}=2.086$ and $z_{\rm spec}=1.953$ from GLASS, respectively. {\it Col. 6}: Demagnified effective radius, $r_{\rm s, demag}$ estimated as $r_{\rm s}/\sqrt{\mu}$, in arcseconds. Upper limits correspond to $2\sigma$. {\it Col. 7}: Demagnified flux density, $F_{\rm uv-fit, demag}$, estimated as $F_{\rm uv-fit}/\mu$, in mJy. {\it Col. 8}: Extended emission classification, with 0 for point-like, 1 for marginally extended, and 2 for significantly extended. \label{tab:gold_size}} \centering \begin{tabular}{cccccccc} \hline {ID} & {$r_{\rm s}$ [$\arcsec$]} & {$q_{\rm s}$} & {$\phi_{\rm s}$ [$^{\circ}$]} & {$\mu$} & $r_{\rm s, demag}$ [$\arcsec$] & {$F_{\rm uv-fit, demag}$ [mJy]} & Extension. \\ \hline \noalign{\smallskip} A2744-ID01 & $ 0.05 \pm 0.01 $ & $ 0.47 \pm 0.21 $ & $ 110 \pm 26 $ & $2.8^{+1.3}_{-0.6}$ & $0.03^{+0.01}_{-0.01}$ & $0.557^{+0.163}_{-0.176}$ & 1 \\[5pt] A2744-ID02 & $ 0.23 \pm 0.04 $ & $ 0.17 \pm 0.05 $ & $ 85 \pm 2 $ & $1.7^{+0.6}_{-0.4}$ & $0.17^{+0.04}_{-0.03}$ & $1.630^{+0.534}_{-0.373}$ & 2 \\[5pt] A2744-ID03 & $ 0.26 \pm 0.03 $ & $ 0.58 \pm 0.13 $ & $ 81 \pm 11 $ & $1.9^{+0.3}_{-0.3}$ & $0.19^{+0.03}_{-0.03}$ & $0.849^{+0.161}_{-0.133}$ & 2 \\[5pt] A2744-ID04& $ 0.06 \pm 0.02 $ & $ 0.62 \pm 0.24 $ & $ 84 \pm 51 $ & $2.7^{+1.2}_{-1.0}$ & $0.04^{+0.02}_{-0.01}$ & $0.372^{+0.197}_{-0.112}$ & 1 \\[5pt] A2744-ID05 & $ 0.19 \pm 0.05 $ & $ 0.66 \pm 0.23 $ & $ 60 \pm 42 $ & $1.6^{+0.3}_{-0.3}$ & $0.16^{+0.04}_{-0.04}$ & $0.685^{+0.197}_{-0.122}$ & 2 \\[5pt] A2744-ID06 & $ 0.26 \pm 0.08 $ & $ 0.3 \pm 0.17 $ & $ 50 \pm 10 $ & $2.2^{+0.7}_{-0.8}$ & $0.18^{+0.06}_{-0.06}$ & $0.577^{+0.285}_{-0.164}$ & 2 \\[5pt] A2744-ID07 & $ 0.07 \pm 0.04 $ & $ 0.56 \pm 0.24 $ & $ 85 \pm 57 $ & $1.6^{+0.2}_{-0.2}$ & $<0.12$ & $0.345^{+0.069}_{-0.059}$ & 0 \\[5pt] MACSJ0416-ID01 & $ 0.23 \pm 0.06 $ & $ 0.61 \pm 0.24 $ & $ 99 \pm 71 $ & $1.8^{+0.1}_{-0.5}$ & $0.18^{+0.05}_{-0.04}$ & $0.773^{+0.239}_{-0.107}$ & 2 \\[5pt] MACSJ0416-ID02 & $ 0.32 \pm 0.15 $ & $ 0.58 \pm 0.23 $ & $ 63 \pm 40 $ & $2.2^{+0.3}_{-0.4}$ & $0.22^{+0.10}_{-0.10}$ & $0.259^{+0.086}_{-0.062}$ & 2 \\[5pt] MACSJ0416-ID03 & $ 0.10 \pm 0.07 $ & $ 0.62 \pm 0.22 $ & $ 97 \pm 51 $ & $1.5^{+0.4}_{-0.4}$ & $<0.20$ & $0.267^{+0.093}_{-0.063}$ & 0 \\[5pt] MACSJ0416-ID04 & $ 0.37 \pm 0.21 $ & $ 0.56 \pm 0.28 $ & $ 81 \pm 62 $ & $1.8^{+0.4}_{-0.5}$ & $0.28^{+0.17}_{-0.16}$ & $0.269^{+0.122}_{-0.096}$ & 1 \\[5pt] MACSJ1149-ID01 & $ 0.28 \pm 0.13 $ & $ 0.61 \pm 0.23 $ & $ 92 \pm 44 $ & $4.2^{+5.4}_{-2.1}$ & $0.13^{+0.10}_{-0.07}$ & $0.137^{+0.148}_{-0.083}$ & 2 \\[5pt] \hline \end{tabular} \end{table*} \begin{figure*}[!htbp] \includegraphics[width=\hsize/2]{plots/point_source_distribution_a2744.pdf} \includegraphics[width=\hsize/2]{plots/point_source_distribution_macs0416.pdf} \caption{Measured effective radii $r_{\rm s}$ as a function of flux density for the ALMA sources. We plot sources from A2744 (\textit{left}) and MACS0416$+$MACS1149 (\textit{right}) separately, since the latter have factors of $\approx$2 larger beam sizes. The solid black line in each plot corresponds to the average value of $r_{\rm s}$ measured for simulated point sources ingested into the data. The dashed black lines likewise correspond to limits of 1, 2 and 3 times the standard deviation ($\sigma$) above the mean for simulated point sources, respectively, as a function of flux density. For visualization purposes, we only plot simulated lines for MACSJ0416 in the right panel, noting that the limits for MACSJ1149 are roughly identical. Additionally, A2744-ID02 is not shown in the left plot, since its flux density is much higher ($\sim$2.8\,mJy) and it is clearly extended. Based on the criteria in $\S$\ref{sec:size}, two sources are considered point-like, three sources are marginally extended, and seven are significantly extended. \label{fig:size_point_source}} \end{figure*} \begin{figure}[!htbp] \centering \includegraphics[width=\hsize]{plots/fluxes_sizes_v2.pdf} \caption{Estimated demagnified sizes, $r_{\rm s, demag}$, for the ALMA-detected sources in the three FFs clusters studied here, as measured by fits in the $uv$-plane, plotted alongside measured sizes for several literature samples obtained using interferometric data \citep{Younger2008, Younger2010, Valtchanov2011, Rybak2015, Simpson2015, Bussmann2015}. The estimated demagnified 1.1\,mm flux densities are shown; for the literature sample, when necessary we have converted to this wavelength assuming a slope of $\beta=1.8$. We see that the ALMA-FFs sources studied here exhibit a large dispersion in extent compared to brighter flux density sources; notably $\sim$50\% of the ALMA-FFs sources appear comparable to the largest literature sample sources. The dashed line shows the best fit regression to all samples, with $\log{r_{\rm s, demag}} = m\times\log{F_{\rm 1.1\,mm}} + b$ with $m=-0.08\pm0.10$ and $b=-0.95\pm0.05$. The orange shaded region shows the $2\sigma$ range of possible fits to the data. The measured sizes show no obvious trend with flux density. \label{fig:fluxes_sizes}} \end{figure} To estimate the sizes of the observed sources (in the natural weighting image plane), we used the 'best-fit' models that were fit in the previous subsection. The source size parameters were estimated from the posterior PDF, in the same manner as for the flux density. The estimated parameters are listed in Table~\ref{tab:gold_size}. The observed effective radii, $r_{\rm s}$, range from $0.05\pm0.01\arcsec$ up to $0.37\pm0.21\arcsec$. The reliability of these measurements is a function of their measured signal-to-noise, as it is harder to measure the emission extent at low $S/N$. In order to verify which sources are truly extended, we ingested fake point sources with a range of flux densities directly into the observed visibility data for each galaxy cluster. The ingested flux densities were chosen to recover similar $S/N$ and flux density values as for the detected sources. Once the sources were ingested, the sizes where measured in the same manner as done for the real detections. Figure \ref{fig:size_point_source} shows the measured effective radii for each detected source as well for the full range of ingested point sources as a function of flux density. As noted previously, the achievable size limit for a point source is less accurate at lower $S/N$, making it more challenging to measure the emission extent. To account for this, we define a source size as significantly extended when its measured effective radius lies $\approx$2$\sigma$ above the mean size measured for point sources with the same flux density (denoted by the second dashed line in Fig. \ref{fig:size_point_source}). Sources with measured extents above this 2$\sigma$ threshold but error bars that drop below it are considered marginally extended, while all remaining sources (i.e., any within $\approx$2$\sigma$ of the mean) are considered point-like. only two sources (A2744-ID07 and MACSJ0416-ID03) are fully consistent with being point sources, while seven are clearly significantly extended. The 2$\sigma$ line for MACSJ1149 is a bit lower than the one plotted for MACSJ0416 so the source is classified as extended despite of its lower error bar going below the 2$\sigma$ line. We consider a further three (A2744-ID01, A2744-ID04 and MACSJ0416-ID04) to be significantly extended, although their error bars put them close to the threshold for the point source size distribution. The extent classifications are listed in the last column of Table~\ref{tab:gold_size}. To test the reliability of the measured sizes for sources with low signal-to-noise which were found to be extended, we simulated 10 extended sources with the same size properties as MACSJ0416-ID04, the detection with the lowest $S/N$ in all three clusters. The assumed properties are $r_{\rm s}=0\farcs37$, $q_{\rm s}=0.56$ and $\phi_{\rm s}=81^{\circ}$. After fitting the sizes of the ingested sources, the average parameter measurements were $r_{\rm s}=0\farcs36\pm0\farcs04$, $q_{\rm s}=0.55\pm0.08$ and $\phi_{\rm s}=83.5^{\circ}\pm13.4^{\circ}$. These all lie within $1\sigma$ of the input parameters for the ingested sources, demonstrating that reliable size measurements can be made for extended sources with $S/N$ values similar to even our weakest detections. Unfortunately, a full characterization of the reliability of size measurements for low signal-to-noise sources using \texttt{uvmcmcfit} over all parameter space is well beyond the scope of this paper, and we refer interested readers to \citet{Bussmann2013,Bussmann2015} for such details. In order to analyze the delensed size distribution of the extended sources, we need to estimate the magnification values of each galaxy behind the galaxy cluster, which relies both on the source redshift and the cluster lensing model. MACSJ0416-ID01 and MACSJ0416-ID02 are the only sources with measured redshifts, $z=2.086$ and $z=1.953$ respectively, obtained from the Grism Lens-Amplified Survey from Space \citep[GLASS, ][]{Treu2015}. For the rest of the sources, we will assume that they lie at $z$$=$2 and use the range $z$$=$1--3 as an estimate of the error associated with this assumption. This average redshift and dispersion are consistent with the redshift distribution of published 1.1\,mm detected galaxies found by ALMA to date \cite{Simpson2014, Dunlop2016}. Since none of the sources are close to the critical curves of the clusters, as seen in Fig. \ref{fig:map_A2744}, \ref{fig:map_MACS0416} and \ref{fig:map_MACS1149}, the redshift uncertainty will not be critical to the magnification estimate. As stated above, each of these cluster fields has a set of lensing models that were created by several independent teams using different assumptions and techniques \citep{Bradac2005,Bradac2009,Diego2015,Jauzac2015a,Jauzac2015b,Jauzac2016,Johnson2014,Kawamata2016,Liesenborgs2006,Merten2009,Merten2011,Sebesta2015,Williams2014,Zitrin2009,Zitrin2013}. We estimate the magnification values as in \cite{Coe2015}, by taking the 16th, 50th and 84th percentiles of the magnification values given by the full range from all the models in the FFs archives. We propagate the $z$$=$1--3 redshift uncertainty for the sources which lack spectroscopic redshifts into the magnification errors. For A2744 and MACS0416 we only use the v3 or newer models, as the v3 models correspond to those made using the deep Frontier Fields images for modeling. For MACSJ1149, no v3 models are available, so we use the latest available model from each team. In Table~\ref{tab:gold_size} we present the derived magnification values for the sources. Most of the magnification values lie around $\mu\sim2$, with the highest reaching a value of $\mu=4.2$. For MACSJ0416-ID01 and MACSJ0416-ID02, the galaxies with spectroscopic redshifts, the magnification errors are produced solely by the systematic errors among the lens models of different teams. A comparison of the magnification errors between sources should demonstrate that our adopted redshift uncertainty is not a dominating factor in the magnification error budget. With magnification estimates in hand, we proceed to determine the intrinsic parameters of the detected sources. We correct the effective radius as $r_{s}/\sqrt{\mu}$ and the flux density as $F_{\rm uv-fit}/\mu$. The obtained values are presented in Table~\ref{tab:gold_size}. The errors in the lensing-corrected values are estimated based on the $1\sigma$ range of the resulting distributions of $r_{s}/\sqrt{\mu}$ and $F_{\rm uv-fit}/\mu$ using the parent distributions. With this information, we now compare our demagnified angular sizes and flux densities with those of bright sources recently constrained by ALMA and previous interferometric studies in Figure~\ref{fig:fluxes_sizes}. The comparison samples correspond to: four bright sources observed with the Submillimeter Array \citep[SMA][]{Ho2004} presented by \cite{Younger2008,Younger2010}; the lensed galaxy SDP.81 observed by \cite{ALMA2015} and with an intrinsic size estimated by \mbox{\cite{Valtchanov2011,Rybak2015}}; objects from the UKIDSS ultra deep survey (UDS) part of the SCUBA-2 cosmology legacy survey and the un-lensed objects from the {\it Herschel} Multi-tiered Extragalactic Survey (HerMES), both observed with ALMA \citep{Simpson2015,Bussmann2015}. We do not include strongly lensed galaxies identified in wide field surveys, since they can affected by the size bias. These type of surveys select galaxies by observed flux density, which is the product of the intrinsic flux density and magnification. This typically favors sources with high magnification values, which in turn biases sources to be preferentially smaller. This effect is produced by the small region of high magnification in the source plane, as small sources near the source plane caustics will have a higher flux-weighted magnification than more extended sources \citep{Serjeant2012,Hezaveh2012,Wardlow2013,Spilker2016}. The ALMA-FFs sample is much less affected by the size bias, since none of our detections lie close to the critical curves and therefore do not have high magnification values. It is immediately clear that the samples studied here probe observed flux densities up to an order of magnitude fainter than those measured in the comparison samples. Roughly 2/3 of the ALMA-FFs sources have effective radii of $r_{\rm s, demag}$$\gtrsim$0\farcs16, compared to the average of $<r_{\rm s}>$$\sim$0\farcs1 measured in brighter samples. This high fraction of more extended sources could imply that the sub-mJy population is intrinsically more extended than the brighter compact sources. However, once we factor in the two point-like sources ($r_{\rm s}$$\lesssim$0\farcs05) and the three marginally extended sources, there appears to be considerable dispersion among the fainter ALMA-FFs population. To investigate this issue further, we fit a linear regression between flux density and effective radius as $\log{r_{\rm s}} = m\times\log{F_{\nu}} + b$ to the ALMA-FFs and HerMES samples, where the same code was used to determine $r_{\rm s}$. For simplicity, we convert the upper limits for the point-like sources in our sample into measured sizes. The best-fit regression values are $m=-0.08\pm0.10$ and $b=-0.95\pm0.05$, demonstrating that there is no obvious size evolution. In Figure~\ref{fig:fluxes_sizes} we show the fit relation (dashed line) together with its $2\sigma$ range (orange region). \subsection{Positional Offsets} \begin{table*} \caption[]{High-significance ($\geq5\sigma$) continuum detections NIR counterpart positions and offsets. {\it Col. 1}: Source ID. {\it Cols. 2-3}: Centroid J2000 position of ID in hh:mm:ss.ss+dd:mm:ss.ss for the NIR counterpart. {\it Cols. 4}: Positional error in arcseconds. {\it Col. 5}: Offset measured in the image plane in arcseconds. {\it Col. 6}: Offset measured in the source plane in arcseconds. \label{tab:counterpart_offset}} \centering \begin{tabular}{cccccc} \hline {ID}& {$\alpha_{\rm J2000}$} & {$\delta_{\rm J2000}$} & $\Delta\alpha$, $\Delta\delta$ &{Image plane offset} & {Source plane offset} \\ & [hh:mm:ss.ss] & [$\pm$dd:mm:ss:ss] & [$\arcsec$] & [$\arcsec$] & [$\arcsec$]\\ \hline \noalign{\smallskip} A2744-ID01 & 00:14:19.8058 & -30:23:07.6094 & 0.002 , 0.002 & $ 0.09 \pm 0.01 $ & $ 0.04 \pm 0.01 $\\ A2744-ID02 & 00:14:18.2000 & -30:24:47.3000 & 0.002 , 0.001 & $ 0.68 \pm 0.05 $ & $ 0.48 \pm 0.04 $\\ A2744-ID03 & 00:14:20.4013 & -30:22:54.6038 & 0.007 , 0.006 & $ 0.18 \pm 0.02 $ & $ 0.16 \pm 0.02 $\\ A2744-ID04 & 00:14:17.5816 & -30:23:00.7324 & 0.003 , 0.004 & $ 0.18 \pm 0.02 $ & $ 0.11 \pm 0.02 $\\ A2744-ID05 & 00:14:19.1273 & -30:22:42.3394 & 0.007 , 0.003 & $ 0.17 \pm 0.04 $ & $ 0.15 \pm 0.04 $\\ A2744-ID06 & 00:14:17.2759 & -30:22:58.8000 & 0.003 , 0.003 & $ 0.22 \pm 0.07 $ & $ 0.12 \pm 0.05 $\\ A2744-ID07 & 00:14:22.1091 & -30:22:49.7821 & 0.004 , 0.002 & $ 0.16 \pm 0.03 $ & $ 0.11 \pm 0.02 $\\ MACSJ0416-ID01 & 04:16:10.7762 & -24:04:47.5002 & 0.001 , 0.001 & $ 0.20 \pm 0.03 $ & $ 0.18 \pm 0.03 $\\ MACSJ0416-ID02 & 04:16:06.9539 & -24:03:59.9277 & 0.002 , 0.003 & $ 0.14 \pm 0.08 $ & $ 0.12 \pm 0.06 $\\ MACSJ0416-ID03 & 04:16:08.8169 & -24:05:22.3970 & 0.003 , 0.003 & $ 0.22 \pm 0.06 $ & $ 0.17 \pm 0.06 $\\ MACSJ0416-ID04 & 04:16:11.6688 & -24:04:19.6271 & 0.002 , 0.002 & $ 0.23 \pm 0.12 $ & $ 0.18 \pm 0.08 $\\ MACSJ1149-ID01 & 11:49:36.0961 & +22:24:24.4659 & 0.002 , 0.001 & $ 0.20 \pm 0.09 $ & $ 0.25 \pm 0.14 $\\ \hline \end{tabular} \end{table*} \begin{figure}[!htbp] \centering \includegraphics[width=\hsize]{plots/offsets.pdf} \caption{Positional offsets between the ALMA detections and the nearest optical/NIR detections. We find good agreement between the positions of the sources in the FIR and in the optical rest-frame. The high-significance continuum detections demonstrate that the astrometric consistency between the {\it HST} and ALMA reference frames appears robust to $\approx$0\farcs1 and that the NIR and FIR/mm emission are nearly co-spatial for most of the cases. \label{fig:offsets}} \end{figure} Because the ALMA-detected sources may be optically faint and/or red, we use the deep {\it HST} WFC3 $F160W$ image in the FFs to search for NIR counterparts. In general, we adopt the closest galaxy as the true counterpart (typical separations smaller than $\sim0.2\arcsec$). Once the counterparts are selected, we fit a simple 2-d elliptical Gaussian to the NIR emission to more accurately measure the positions of the counterparts centroids in the same manner as the positions of the ALMA sources. In Figure~\ref{fig:offsets}, we present the measured positional offsets between the 1.1\,mm and NIR emission of the detected galaxies. A2744-ID02 was left out of this analysis, as it appears to arise from a very obscured part of an extended counterpart galaxy, as described in $\S$\ref{Sec:high_significance_cont_detections}. This offset appears real and highlights an important case when the FIR/mm and optical/NIR are not co-spatial, similar to physical offsets seen in the local starburst galaxy NGC4038/9 \citep{Wang2004,Klaas2010}. The intrinsic offset between the FIR/mm and optical/NIR emission can be enlarged by lensing when observed in the image plane, as in our case. Since one of the goals of this analysis is to check the astrometric agreement between {\it HST} and ALMA, we therefore exclude this galaxy from our offset measurement. The measured median offsets between the NIR and mm sources are $\Delta{\rm RA}$$=$0\farcs02$\pm$0\farcs03 and $\Delta{\rm DEC}$$=$$-$0\farcs13$\pm$0\farcs02. with a combined offset scatter of $\approx0\farcs1$. The errors in the offsets include the uncertainties of the NIR and ALMA emissions centers presented in Table \ref{tab:counterpart_offset} and plotted in Figure \ref{fig:counterparts}. These offsets are consistent with the astrometric offsets ($\approx0\farcs1$) measured for the Subaru catalogs to which the {\it HST} A2744 data was tied. \footnote{\url{https://archive.stsci.edu/pub/hlsp/frontier/abell2744/catalogs/subaru/astrometry/}}. By comparing the observed positions of the phase calibrators used in the ALMA observations to those in the literature we find offsets of $\Delta{\rm RA}$=$<0\farcs01$ and $\Delta{\rm DEC}$=$<0\farcs01$ for A2744 and MACSJ1149 and $\Delta{\rm RA}$=$0\farcs012$ and $\Delta{\rm DEC}$=$<0\farcs016$ for MACSJ0416. Although the measured offset in MACSJ0416 is somewhat larger than those measured in the other two clusters, all are minimal compared to the observed offsets in the detected sources, indicating that the measured offsets are not due to astrometric problems in the ALMA observations. The small scatter measured for the secure sample indicates that the NIR and FIR/mm emission are nearly co-spatial, except for the the special case of A2744-ID02 described above. These results also demonstrate that the astrometric consistency between the {\it HST} and ALMA reference frames appears robust to $\approx$0\farcs1. The measured offsets could be due to small astrometric offsets or true spatial offsets between the sources optical/NIR and FIR/mm emission. Under the assumption that we have a perfect astrometric agreement between ALMA and {\it HST}, the average offset for all the sources plotted in Fig.~\ref{fig:offsets} is $0\farcs17\pm0\farcs02$. To rule out the effects of lensing, we tracked the positions of the centroids in ALMA and {\it HST} to the source plane using the lensing models of Zitrin-NFWv3 and Zitrin-LTMv1 and measured intrinsic offsets. We selected Zitrin-NFW models since they were observed to better follow the median distribution of magnifications given by the combinations of all lens models \citet{Priewe2016}. This measurement includes the effect of lensing shear. Individual source plane offsets are given in Table~\ref{tab:counterpart_offset}, while the average demagnified offset is still $0\farcs14\pm0\farcs02$, which is $\approx$1.2\,kpc at $z=2$. The source plane offset errors include the positional uncertainties in both the NIR and ALMA emission but we have not propagated the systematic lens models uncertainties. For A2744-ID02, the demagnified offset to the brightest peak, assuming $z=2$, is $0\farcs48\pm0\farcs04$ including only the uncertainties in the position of both NIR and ALMA emissions. This offsets corresponds to $4.1\pm0.3$ kpc at $z=2$ similar than the spatial offset of $\sim4$ kpc measured in GN20 at $z=4.05$, one of the best studied SMGs at high redshift \citep{Hodge2015}. It is clear that these ALMA observations can open a window to study and resolve the obscured star-formation activity in galaxies at high redshift. While A2744-ID02 is the most obvious case, an interesting aspect of the counterpart offsets is that for nine out of 12 ALMA-FFs sources (i.e., A2744-ID01, A2744-ID02, A2744-ID03, A2744-ID04, A2744-ID06, A2744-ID07, M0416-ID02, M0416-ID03, and M0416-ID04), the ALMA centroid position falls on a darker portion of the counterpart galaxy compared to the NIR peak (see Figure~\ref{fig:cont_fits1}). This effect has been noted other ALMA surveys \citep[e.g.,][]{Wiklind2014, Dunlop2016}, and attributed to the fact that the FIR/mm emission region likely suffers from strong dust extinction. This physical effect may be an important term in the remaining measured offsets, although the current error bars make this difficult to constrain for individual sources.
16
7
1607.03808
Context. Dusty star-forming galaxies are among the most prodigious systems at high redshift (z &gt; 1), characterized by high star-formation rates and huge dust reservoirs. The bright end of this population has been well characterized in recent years, but considerable uncertainties remain for fainter dusty star-forming galaxies, which are responsible for the bulk of star formation at high redshift and thus play a key role in galaxy growth and evolution. <BR /> Aims: In this first paper of our series, we describe our methods for finding high redshift faint dusty galaxies using millimeter observations with ALMA. <BR /> Methods: We obtained ALMA 1.1 mm mosaic images for three strong-lensing galaxy clusters from the Frontier Fields Survey, which constitute some of the best studied gravitational lenses to date. The ≈2' × 2' mosaics overlap with the deep HST WFC3/IR footprints and encompass the high magnification regions of each cluster for maximum intrinsic source sensitivity. The combination of extremely high ALMA sensitivity and the magnification power of these clusters allows us to systematically probe the sub-mJy population of dusty star-forming galaxies over a large surveyed area. <BR /> Results: We present a description of the reduction and analysis of the ALMA continuum observations for the galaxy clusters Abell 2744 (z = 0.308), MACS J0416.1-2403 (z = 0.396) and MACS J1149.5+2223 (z = 0.543), for which we reach observed rms sensitivities of 55, 59 and 71 μJy beam<SUP>-1</SUP> respectively. We detect 12 dusty star-forming galaxies at S/N ≥ 5.0 across the three clusters, all of them presenting coincidence with near-infrared detected counterparts in the HST images. None of the sources fall close to the lensing caustics, thus they are not strongly lensed. The observed 1.1 mm flux densities for the total sample of galaxies range from 0.41 to 2.82 mJy, with observed effective radii spanning ≲0.̋05 to 0.̋37 ± 0.̋21 . The lensing-corrected sizes of the detected sources appear to be in the same range as those measured in brighter samples, albeit with possibly larger dispersion. <P />Reduced mosaic images are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (<A href="http://130.79.128.5">http://130.79.128.5</A>) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/597/A41">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/597/A41</A>
false
[ "high redshift faint dusty galaxies", "galaxy growth", "galaxies", "high redshift", "star formation", "maximum intrinsic source sensitivity", "fainter dusty star-forming galaxies", "dusty star-forming galaxies", "MACS", "Dusty star-forming galaxies", "observed rms sensitivities", "huge dust reservoirs", "anonymous ftp", "the galaxy clusters", "high star-formation rates", "observed effective radii", "millimeter observations", "date", "brighter samples", "±" ]
13.236951
7.142813
132
1991062
[ "Yoon, Jinmi", "Beers, Timothy C.", "Placco, Vinicius M.", "Rasmussen, Kaitlin C.", "Carollo, Daniela", "He, Siyu", "Hansen, Terese T.", "Roederer, Ian U.", "Zeanah, Jeff" ]
2016ApJ...833...20Y
[ "Observational Constraints on First-star Nucleosynthesis. I. Evidence for Multiple Progenitors of CEMP-No Stars" ]
156
[ "Department of Physics, University of Notre Dame, Notre Dame, IN 46556, USA ; Joint Institute for Nuclear Astrophysics-Center for the Evolution of the Elements (JINA-CEE), USA;", "Department of Physics, University of Notre Dame, Notre Dame, IN 46556, USA ; Joint Institute for Nuclear Astrophysics-Center for the Evolution of the Elements (JINA-CEE), USA", "Department of Physics, University of Notre Dame, Notre Dame, IN 46556, USA ; Joint Institute for Nuclear Astrophysics-Center for the Evolution of the Elements (JINA-CEE), USA", "Department of Physics, University of Notre Dame, Notre Dame, IN 46556, USA ; Joint Institute for Nuclear Astrophysics-Center for the Evolution of the Elements (JINA-CEE), USA", "Department of Physics, University of Notre Dame, Notre Dame, IN 46556, USA ; Joint Institute for Nuclear Astrophysics-Center for the Evolution of the Elements (JINA-CEE), USA", "Department of Physics, Xi'an Jiaotong University, Shaanxi, 710049, People's Republic of China", "Observatories of the Carnegie Institution of Washington, 813 Santa Barbara Street, Pasadena, CA 91101, USA", "Joint Institute for Nuclear Astrophysics-Center for the Evolution of the Elements (JINA-CEE), USA ; Department of Astronomy, University of Michigan, 1085 S. University Avenue, Ann Arbor, MI 48109, USA", "Z Solutions, Inc., 9430 Huntcliff Trace, Atlanta, GA 30350, USA" ]
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[ "astronomy" ]
11
[ "stars: abundances", "stars: AGB and post-AGB", "stars: chemically peculiar", "stars: evolution", "stars: massive", "stars: Population II", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.3847/0004-637X/833/1/20", "10.48550/arXiv.1607.06336" ]
1607
1607.06336_arXiv.txt
At low iron abundances relative to the Sun, a substantial fraction of the stars in the halo of the Milky Way have been found to be greatly enhanced in carbon, the so-called carbon-enhanced metal-poor (CEMP) stars. \cite{beers2005} originally divided such stars into several sub-classes, depending on the nature of their neutron-capture element abundance ratios -- CEMP-$s$, CEMP-$r$, CEMP-$r/s$, and CEMP-no\footnote{CEMP-$s$ : [C/Fe] $>$ +1.0, [Ba/Fe] $>$ +1.0, and [Ba/Eu] $>$ +0.5\newline CEMP-$r$ : [C/Fe] $>$ +1.0 and [Eu/Fe]$>$ +1.0 \newline CEMP-$r/s$ : [C/Fe] $>$ +1.0 and 0.0 $<$ [Ba/Eu] $<$+0.5 \newline CEMP-no : [C/Fe] $>$ +1.0 and [Ba/Fe] $<$ 0.0}. As discussed by these authors, and many since, the observed differences in the chemical signatures of the sub-classes of CEMP stars are thought to be tied to differences in the astrophysical sites responsible for the nucleosynthesis products they now incorporate in their atmospheres, including elements produced by the very first generations of stars. \subsection{The Origin of CEMP Stars} In this paper we focus on the two most populous sub-classes of the carbon-enhanced metal-poor stars, the CEMP-$s$ and CEMP-no stars. Based on both extensive observational follow-up and theoretical modeling, the elemental abundance pattern associated with the CEMP-$s$ stars (carbon enhancement accompanied by strong over-abundances of neutron-capture elements produced by the main $s$-process) is thought to arise from an {\it extrinsic} process -- mass transfer to the presently observed star from an evolved binary companion. This companion, the site where the enhancement of carbon and the $s$-process elements originally took place, is expected to have been a low- to intermediate-mass ($\sim$1 to $\sim$4 $M_{\odot}$) asymptotic giant-branch (AGB) star, which has now evolved to become a faint white dwarf \citep[e.g.,][]{suda2004, herwig2005, lucatello2005, bisterzo2011,starkenburg2014,hansen2015a}. The mass-transfer process itself has proven challenging to model, despite extensive efforts in recent years (see, e.g., \citealt{abate2013, abate2015c,abate2015b}, and references therein). Accreted material can potentially be mixed into the atmosphere of the presently observed companion by several processes (e.g., thermohaline mixing, levitation, etc., see \citealt{stancliffe2007, stancliffe2008,stancliffe2010,matrozis2016}). Additional processing can also occur towards the tip of the giant branch \citep{placco2013, karakas2014,placco2014c}, all of which may complicate interpretation of the observed elemental-abundance patterns. Binary mass transfer is thought to play a role in the origin of the CEMP-$r/s$ stars as well \citep[e.g.,][]{jonsell2006,lugaro2009, bisterzo2011,herwig2011}, but the origin of the $r$-process-like component of their abundance pattern remains unclear \citep{abate2016}. It remains possible that yet another nucleosynthetic process, the so-called $i$-process, may need to be invoked to account for their observed abundance patterns \citep{dardelet2015,hampel2016}. For the purpose of the present analysis, we group the CEMP-$r/s$ stars along with the CEMP-$s$ stars. In contrast to the CEMP-$s$ and CEMP-$r/s$ stars, a number of lines of observational evidence (including long-term radial-velocity monitoring; see \citealt{hansen2016a}) indicate that the distinctive abundance patterns of CEMP-no stars (carbon enhancement with a lack of neutron-capture element over-abundances) arose from an {\it intrinsic} process\footnote{We intend this term to indicate that the observed elemental abundances on the surface of the star were present in the gas from which the star first formed, and not (as it is also used) patterns arising from internal processing in the star of material that is later transported to the stellar surface.}. The inference is that the presently observed CEMP-no stars are indeed bona-fide second-generation stars, born in natal clouds polluted by massive first-generation stars. A number of astrophysical sites for the progenitors of the CEMP-no stars have been suggested. The so-called ``faint supernovae'' or ``mixing-and-fallback'' models \citep{umeda2003,umeda2005,nomoto2013, tominaga2014} hold that the gas from which CEMP-no stars formed was enriched by a supernova without sufficient explosion energy to release its full complement of synthesized heavy elements (which fall back to the nascent neutron star or black hole at its center), and only the lighter elements (including C, N, O, and other light elements such as Na, Mg, Al, and Si) are expelled. Another possibility, the so-called ``spinstar'' model \citep[e.g.,][]{meynet2006,meynet2010,chiappini2013} proposes that the gas from which CEMP-no stars formed was enhanced in carbon (as well as N and O) by the strong stellar winds expected to arise from rapidly-rotating massive stars of ultra low metallicity. In addition, \citet{heger2010} have considered possible progenitors of the CEMP-no stars including the effects of rotation and mixing and fallback. More recent modeling has suggested that spinstars may also be capable of producing other light elements and some amount of first-peak neutron-capture elements (such as Sr) and second-peak $s$-process elements (such as Ba), and possibly even third-peak elements such as Pb, depending on the degree of the internal mixing induced by the rapid rotation~\citep{maeder2015,frischknecht2016}. Finally, we note that \citet{cooke2011a,cooke2012} have reported on recently discovered high-redshift carbon-enhanced damped Lyman-$\alpha$ systems that exhibit elemental-abundance patterns which resemble those expected from massive, carbon-producing first stars. These authors speculated that the progenitors that produced these patterns are the same as those responsible for those associated with CEMP-no stars in the Galaxy. \subsection{The High and Low Carbon Bands for CEMP Stars} \citet{spite2013} used literature abundance data for $\sim 50$ CEMP main-sequence turnoff and dwarf stars, including both CEMP-$s$ and CEMP-no stars, and plotted the absolute carbon abundance, $A$(C) $= \log\,\epsilon$(C)\footnote{$A$(X) = $\log\,\epsilon$(X) = $\log\,$($N_{\rm X}/N_{\rm H}$)+12, where X represents a given element.}, as a function of metallicity, [Fe/H], for their sample. The stars in their sample were specifically chosen to be in evolutionary stages where alteration of their surface elemental abundances, due to significant internal mixing, were not expected to have occurred. We point out, however, that certain processes, such as thermohaline mixing, occur almost immediately after mass-transfer events (in CEMP-$s$ stars; R. Stancliffe, priv. comm.), so some mixing (dilution) may have occured even in supposedly unevolved stars. Based on this sample, they claimed the existence of a clear bimodality among the CEMP stars -- the stars in their sample with [Fe/H] $> -$3.0, which are dominated by CEMP-$s$ stars, populate a high-C ``plateau" at $A$(C) $\sim 8.25$, close to the Solar value of $A$(C). In contrast, the stars with [Fe/H] $ < -$3.4, which are exclusively CEMP-no stars, reside in a lower region (and possible plateau) at $A$(C) $\sim 6.5$. They interpreted this behavior as the result of the different carbon-production mechanisms for these sub-classes of stars -- mass-transfer from binary AGB companions in the case of CEMP-$s$ stars and enrichment of the natal clouds of the CEMP-no stars by massive-star nucleosynthesis. \citet{bonifacio2015} confirmed and extended the claim by Spite et al. with a larger sample ($\sim$70) of unevolved main-sequence turnoff and dwarf stars, along with a few lower red giant-branch (RGB) CEMP stars with [Fe/H] $ > -3.5$. These authors found a clear separation of the $A$(C) distribution, but commented that the individual distributions of carbon abundance were quite wide, on the order of one dex. They advocated for a similar explanation of this separation as in Spite et al., based on different carbon-production mechanisms for the CEMP-$s$ and CEMP-no stars. The work of \citet{hansen2015a} provided new data for additional CEMP stars, and considered them along with literature data (compilation from \citealt{yong2013}), confirming once again the existence of the carbon bands, based on a total of 64 stars. However, they identified three CEMP-no stars located on the high-C band, as well as the apparent existence of a smooth transition of $A$(C) between the two bands, which as they noted presents a challenge to the interpretation of the bimodality in $A$(C) as exclusively due to extrinsic (AGB mass-transfer) and intrinsic (C-enriched ISM) processes. These authors emphasized the crucial role that knowledge of the binary status for stars associated with the two carbon bands may play for determination of the nature of their progenitors, and recently published the results of long-term radial-velocity monitoring for samples of CEMP-no \citep{hansen2016a} and CEMP-$s$ \citep{hansen2016b} stars. In order to further explore these questions, we have compiled an extensive set of 305 CEMP stars with available high-resolution spectroscopic data from the literature, including more recent data than was available to the studies conducted in the past few years. This paper is arranged as follows. Section~\ref{data} describes details of the literature data compilation, and the corrections we have applied in order to place the data on a suitable common scale. The results of our analysis, presented in Section~\ref{results}, clearly support the existence of (at least) a bimodality in the distribution of $A$(C) for CEMP stars, but we note that the $A$(C) distribution exhibits more complex behavior that is {\it not captured} by its description as carbon plateaus or bands. Instead, we suggest that the CEMP stars can be more usefully described as falling into three groups, one for the CEMP-$s$ (and CEMP-$r/s$) stars and two for the CEMP-no stars, based on their location in the $A$(C)-[Fe/H] space. We discuss these divisions in more detail in Section~\ref{discussion}, and demonstrate the existence of a correlated behavior between the absolute abundances of the light elements Na and Mg, $A$(Na) and $A$(Mg), with $A$(C). Collectively, this may provide the first evidence for the existence of at least two distinct progenitor populations that are responsible for the abundance signatures among CEMP-no stars. In this section we also consider information that can be gleaned from the subset of CEMP stars with known binary status, concluding that the carbon enhancement of the great majority of CEMP-no stars is an intrinsic process, while most CEMP-$s$ (and CEMP-$r/s$) stars are extrinsically enriched, as previously suggested. We also identify several interesting subsets of stars that exhibit abundance anomalies relative to the majority of other CEMP stars in our sample. Finally, we argue that classification based on $A$(C) is likely to be a more astrophysically fundamental (and equally successful) method to distinguish the CEMP-no stars from the CEMP-$s$ and CEMP-$r/s$ stars than the previously employed approach based on [Ba/Fe] (and [Ba/Eu]) ratios, with the considerable advantage that it can be obtained from low- to medium-resolution, rather than high-resolution, spectroscopy. Our conclusions are briefly summarized in Section \ref{conclusion}.
We have investigated the absolute carbon-abundance distribution of CEMP stars based on a reasonably complete compilation of available high-resolution spectroscopic data. The $A$(C) distribution of the CEMP stars clearly exhibits (at least) a bimodality, as has been noted by a number of previous authors. However, there exist complex behaviors embedded in the $A$(C)-[Fe/H] space, not easily captured by description as plateaus or bands; we suggest use of the terms high-C and low-C regions. We separate CEMP stars into three groups -- Group~I, comprising primarily CEMP-$s/rs$ stars, and Group~II and Group~III comprising CEMP-no stars. Along with the apparent dichotomy in the absolute abundance distribution of Na and Mg as a function of $A$(C) for the CEMP-no stars, we suggest this provides the first clear observational evidence for the existence of multiple progenitor populations of the CEMP-no stars in the early Universe. Based on the known binary status for a subset of the CEMP stars, we strongly support the hypothesis that the carbon enhancement of the CEMP-$s/rs$ stars is extrinsic, and due to mass-transfer of material enriched by an AGB companion, while the carbon enhancement of CEMP-no stars is intrinsic, and due to enrichment of their natal clouds by high-mass progenitor stars. According to this view, the CEMP-no stars are bona-fide second generation stars, as supported by other lines of evidence (see summary in \citealt{hansen2016a}). We have identified a number of interesting outliers, worthy of further exploration, that differ from the general behaviors of otherwise similar stars in either their binary status or with disparate $A$(C) and [Ba/Fe] ratios. Finally, we have presented evidence that the separation of CEMP-$s/rs$ stars from CEMP-no stars can be accomplished as well (or better) using the simple criterion $A$(C) $> 7.1$ for CEMP-$s/rs$ stars and $A$(C) $\leq 7.1$ for CEMP-no stars. As $A$(C) can be obtained from low- to medium-resolution spectroscopy, rather than the high-resolution spectroscopy required for the former criterion based on [Ba/Fe] (and [Ba/Eu]) ratios, this provides an efficient means to quickly isolate the most interesting CEMP sub-classes in massive spectroscopic surveys now and in the future. Given the multiple nucleosynthetic pathways for the production of Ba in the early Universe known (or suggested) to exist, we assert that the $A$(C) criterion may also be more astrophysically fundamental. We plan to employ the $A$(C) sub-classification approach to the thousands of CEMP stars presently identified by SDSS/SEGUE and other large surveys from the past, e.g., the HK survey \citep{beers1985, beers1992} and the Hamburg/ESO survey \citep{christlieb2003}, enabling consideration of potential differences in the spatial and kinematic distributions of CEMP-$s/rs$ and CEMP-no stars (e.g., \citealt{carollo2014}), as well as in their relative frequencies as a function of [Fe/H] (e.g., \citealt{lee2013}). Comparison of these observables with the predictions of modern Galactic chemical evolution models (e.g., \citealt{cote2016,crosby2016}) should prove illuminating. Future progress requires a significant increase in the numbers of CEMP-$s/rs$ stars and CEMP-no stars with available high-resolution spectroscopy, so that their full elemental-abundance distributions can be considered in more detail (both from the ground and in the near-UV from space), and used to better constrain their likely progenitors. This goal is being actively pursued. Additional long-term radial-velocity monitoring of CEMP stars is also likely to pay substantial scientific dividends. For both reasons, the identification of, in particular, bright CEMP stars is being given high priority in our ongoing survey efforts. There is also a clear need for additional development of theory and modeling, to obtain deeper understanding of the nucleosynthesis processes in operation in the suggested progenitors of CEMP stars. Links between the CEMP phenomenon and the early-Universe initial mass function have been considered by a number of previous authors (e.g., \citealt{lucatello2005b,tumlinson2007b,tumlinson2007a,suda2013,carollo2014,lee2014}), based on the relative frequencies of CEMP-$s/rs$ and CEMP-no stars in the halo system. Here we have shown that this idea might be extended by taking into account the presumed differences in the masses of the progenitors of the Group~II and Group~III CEMP-no stars.
16
7
1607.06336
We investigate anew the distribution of absolute carbon abundance, A(C) = log ɛ(C), for carbon-enhanced metal-poor (CEMP) stars in the halo of the Milky Way, based on high-resolution spectroscopic data for a total sample of 305 CEMP stars. The sample includes 147 CEMP-s (and CEMP-r/s) stars, 127 CEMP-no stars, and 31 CEMP stars that are unclassified, based on the currently employed [Ba/Fe] criterion. We confirm previous claims that the distribution of A(C) for CEMP stars is (at least) bimodal, with newly determined peaks centered on A(C) = 7.96 (the high-C region) and A(C) = 6.28 (the low-C region). A very high fraction of CEMP-s (and CEMP-r/s) stars belongs to the high-C region, while the great majority of CEMP-no stars resides in the low-C region. However, there exists complexity in the morphology of the A(C)-[Fe/H] space for the CEMP-no stars, a first indication that more than one class of first-generation stellar progenitors may be required to account for their observed abundances. The two groups of CEMP-no stars we identify exhibit clearly different locations in the A(Na)-A(C) and A(Mg)-A(C) spaces, also suggesting multiple progenitors. The clear distinction in A(C) between the CEMP-s (and CEMP-r/s) stars and the CEMP-no stars appears to be as successful, and likely more astrophysically fundamental, for the separation of these sub-classes as the previously recommended criterion based on [Ba/Fe] (and [Ba/Eu]) abundance ratios. This result opens the window for its application to present and future large-scale low- and medium-resolution spectroscopic surveys.
false
[ "CEMP stars", "CEMP", "absolute carbon abundance", "305 CEMP stars", "31 CEMP stars", "multiple progenitors", "A(C", "s", "their observed abundances", "classes", "high-resolution spectroscopic data", "r", "first-generation stellar progenitors", "different locations", "first", "Ba/Fe", "the CEMP-no stars", "-", "A(Na)-A(C", "future large-scale low- and medium-resolution spectroscopic surveys" ]
8.10605
9.06214
135
12510013
[ "Battye, Richard A.", "Pace, Francesco" ]
2016PhRvD..94f3513B
[ "Approximation of the potential in scalar field dark energy models" ]
24
[ "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, United Kingdom", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, United Kingdom" ]
[ "2017CaJPh..95.1068R", "2017JCAP...10..040P", "2017MNRAS.466.1839P", "2017PhRvD..96f4041B", "2018JCAP...01..018S", "2018PhRvD..97j4070B", "2019Ap&SS.364..140P", "2019JApA...40...44R", "2019JCAP...09..060P", "2020CaJPh..98.1119R", "2020IJGMM..1750139A", "2020MPLA...3550166P", "2020Univ....6..244C", "2021CaJPh..99..991M", "2021ZNatA..76...43C", "2022ChPhC..46l5104G", "2022MPLA...3750228M", "2022Univ....8..145P", "2022Univ....8..520C", "2023EPJC...83..918C", "2023JCAP...10..036L", "2023PDU....3901167M", "2024NewA..10902203S", "2024PDU....4501511L" ]
[ "astronomy", "physics" ]
2
[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
[ "1987PhRvD..35.2339F", "1988ApJ...325L..17P", "1988NuPhB.302..668W", "1988PhRvD..37.3406R", "1991CQGra...8..667E", "1993PhRvD..48.2529C", "1998PhRvD..57.4686C", "1998PhRvL..80.1582C", "1999PhLB..458..209A", "1999PhLB..458..219G", "1999PhRvD..59l3504S", "2000IJMPD...9..373S", "2000PhRvD..61h1301G", "2000PhRvD..61l7301B", "2000PhRvD..62b3511C", "2000PhRvL..85.4438A", "2001IJMPD..10..213C", "2001PhRvD..63j3510A", "2002JHEP...04..048S", "2002MPLA...17.1797S", "2002PhLB..545...23C", "2002PhRvD..66b1301P", "2003ApJ...597....9D", "2003JETPL..77..201S", "2003MNRAS.344.1057A", "2003PhLB..575..165A", "2003PhR...373....1G", "2003PhRvL..90i1301L", "2004ApJ...612..652D", "2004JCAP...07..004P", "2004JCAP...09..001H", "2004JHEP...05..074H", "2004PhRvD..69l3502A", "2004PhRvD..69l3517C", "2004PhRvD..70j3505S", "2004PhRvD..70l3505A", "2004PhRvL..93a1301S", "2005IJMPA..20.5513S", "2005PhRvD..71d3003C", "2005PhRvD..71l3001S", "2005PhRvD..72h3511S", "2005PhRvD..72j3503J", "2006CQGra..23.1557R", "2006IJMPD..15.2105S", "2006JCAP...02..004M", "2007PhRvL..98y1301C", "2008ChPhL..25..344Y", "2008JCAP...04..035G", "2008PhLB..666..415B", "2008PhRvD..78j3502S", "2008PhRvL.101r1301Z", "2009JCAP...12..025C", "2009PhRvD..79h3517C", "2009PhRvD..80d3517C", "2009PhRvD..80j3508B", "2009PhRvD..80j9902C", "2011PhRvD..83f3502D", "2011arXiv1110.3193L", "2012EPJC...72.2067S", "2012EPJC...72.2095Y", "2012MNRAS.422.1186P", "2012arXiv1211.0310L", "2013LRR....16....6A", "2013PhRvD..87b3502D", "2014JPhCS.485a2017D", "2014MNRAS.438.1948S", "2015Ap&SS.356..399Y", "2015EPJC...75..244M", "2015JCAP...04..048B", "2015PhRvD..91l3535P", "2015aska.confE..19S", "2015aska.confE..24B", "2015aska.confE..25C", "2015aska.confE..31R", "2016A&A...594A..13P", "2016A&A...594A..14P" ]
[ "10.1103/PhysRevD.94.063513", "10.48550/arXiv.1607.01720" ]
1607
1607.01720_arXiv.txt
The origin of the cosmic acceleration is one of the most significant open questions in cosmology and fundamental physics. A cosmological constant is still very much consistent with the data \citep{Planck2015_XIII,Planck2015_XIV}, but in order to either refute or confirm this simple hypothesis one needs to consider alternative models to explain the observations. One very simple idea is to postulate a dark energy component dominated by a scalar field either with a canonical or noncanonical kinetic term. Such models are known as quintessence models \citep{Ford1987,Peebles1988,Ratra1988a,Wetterich1988,Caldwell1998,Copeland1998,Steinhardt1999,Barreiro2000} and $k$-essence models \citep{ArmendarizPicon1999,Chiba2000,Mukhanov2006}, respectively. The standard approach when constraining cosmological models with a dark energy component that is not the cosmological constant is to define an equation-of-state parameter $w_{\phi}=P_{\phi}/\rho_{\phi}\neq -1$, where $P_{\phi}$ is the pressure of dark energy and $\rho_{\phi}$ is its density, making no assumption as to the origin of the dark energy. In principle this is a general function of time, but it is often considered to be either constant, or to be represented by a specific functional form, for example \cite{Chevallier2001,Linder2003}. At the moment the data barely constrain anything beyond a constant $w_{\phi}$, but this is likely to change in the near future as more observations probing the equation of state become available, such as Euclid\footnote{\url{http://www.euclid-ec.org/}} \citep{Laureijs2011,Amendola2013}, LSST\footnote{\url{http://www.lsst.org}} \citep{LSST2012} and SKA\footnote{\url{https://www.skatelescope.org/}} \citep{Bull2015,Camera2015,Raccanelli2015,Santos2015}. Various ideas have been put forward to extend to time varying situations. These include various limited functional forms \citep{Chevallier2001,Linder2003,Hannestad2004,Jassal2005,Barboza2008}, the \textit{Om} diagnostic \citep{Sahni2008,Zunckel2008}, the state-finder approach \citep{Alam2003,Sahni2003} and even using principal component analysis on general piecewise linear parametrizations of $w_{\phi}$ \citep{Crittenden2009}. For a review of the parametric and nonparametric methods to reconstruct the dark energy equation-of-state parameter, we refer to \cite{Sahni2006}. Since many of the observations are sensitive to perturbations in the dark energy it is also necessary to make some assumptions about the perturbations, but we will not consider this here. An alternative is to presume that the origin of the dark energy is a model based on a scalar field. However, such models usually involve one or more arbitrary functions which would need to be specified before any model prediction could be made. One of these is the potential $V(\phi)$ of the scalar field which one might try to reconstruct from observations. One obvious suggestion \citep{Sahlen2005}, which extends the approach of \cite{Grivell2000} for inflation, is to represent the potential as a Taylor series expanded around the present-day value of the field $\phi_0$ \begin{equation}\label{eqn:Vphi_series} V(\phi)=V_0+V_1(\phi-\phi_0)+V_2(\phi-\phi_0)^2+\dots\;, \end{equation} and attempt to fit for the coefficients $V_i$. However, it is not clear where to truncate this series in a controlled way. Similar and complementary methods have been proposed by \cite{Copeland1993,Daly2003,Daly2004,Simon2005a}. Other reconstruction methods are valid in the slow-roll regime, that is, when $1+w_{\phi}\approx 0$. For quintessence models, a one-parameter \citep{Slepian2014} or two-parameter \citep{Crittenden2007,Chiba2009a,Chiba2009b} formula has been used and for $k$-essence models we refer to works by \cite{ArmendarizPicon1999,Chiba2009c}. In this paper we first calculate potentials for a range of minimally coupled scalar field models with canonical (\autoref{sect:mcsf}) and noncanonical (\autoref{sect:kessence}) kinetic terms assuming initially that $w_{\phi}$ is constant. It is possible to derive an analytic solution for the potential in Quintessence models, but this is not possible in general for the case of $k$-essence models and therefore we resort to numerical calculations and an analytic approximation around the present day which is valid for $0.6\lesssim a\lesssim 1.4$. Based on this analytic approximation we suggest a form of a potential with just four parameters which we demonstrate can lead to a wide range of behaviour for $w_{\phi}$ as a function of time (\autoref{sect:wa}) and, by design, includes models with constant $w_{\phi}$. Of course, this functional form will not include every possible behaviour in a general model, but it does provide more physical insights and it is useful for models which are not significantly different from a linearly evolving equation-of-state parameter. We conclude and discuss our results in \autoref{sect:conclusions}. In the following, we will use natural units with $c=\hbar=1$, the Planck mass is $M_{\rm pl}=G^{-1/2}$ and we assume a metric with signature $(-,+,+,+)$.
\label{sect:conclusions} Scalar fields are an important field of research in cosmology and are one of the most studied candidates used to explain and describe the accelerated expansion of the Universe. In this work, we consider two main classes of models: minimally coupled models (both quintessence and phantom) and $k$-essence models. For this second class, we specialise the Lagrangian to assume three particular functional forms, dubbed type A, type B and type C models. In each case, we have shown that specifying the scalar field potential $V(\phi)$, one can determine the evolution of the scalar field and the corresponding equation of state $w_{\phi}(a)$. This is true generally but in order to make it clear, we have assumed the equation of state to be known and we calculated explicitly the time evolution of the scalar field and of the potential in some cases. We showed that it is possible to obtain an exact analytic solution for minimally coupled and the type C models with constant equation of state. This is not possible for more general $k$-essence models or for models with a time-varying $w_{\phi}(a)$, but we have solutions for $\phi(a)$ as definite integrals and these can be used to establish the potentials, $V(\phi)$ numerically. We have also derived useful approximate forms of the potential which are valid in different epochs, corresponding to the domination of one cosmic fluid. In particular we deduce the form of the potential at early times ($a\ll 1$, corresponding to the matter dominated epoch) and at late times ($a\gg 1$, corresponding to the scalar field dominated regime), showing that in general the potential is often very well approximated by a power-law. From an observational point of view, the most important regime to understand the potential is around $a\approx 1$. Assuming initially a constant equation of state $w_{\phi}$, we showed that the scalar field potential can be approximated by the expression given in (\ref{eqn:Vphi_ABCD}). This expression depends only on four parameters and with the appropriate choice of coefficients can cover all the classes of models studied in this work. In \autoref{sect:wa} we discussed how this expression might be applied to dynamical dark energy models, by appropriately choosing a new set of parameters which reduces to the correct expression in the limit of constant $w_{\phi}$. Note that this can not be done for type B models, since our formalism only works for constant equations of state. In some respect our approach is similar to the work of \cite{Sahlen2005}. To derive our expression in (\ref{eqn:Vphi_ABCD}), we performed a Taylor expansion of the scalar field evolution, so the same critique could be applied: where to stop the series? Our approximate potential, for $a\approx 1$ ($\phi\approx\phi_0$) can be expanded in powers of $\phi-\phi_0$ leading to the same form of the potential proposed by \cite{Sahlen2005}. In contrast to that work, our proposed potential has well motivated coefficients and in the regime of interest it would be possible to map the $V_i$ coefficients of (\ref{eqn:Vphi_series}) in terms of our four parameters. For example, at zeroth order, we can write $V_0$ in (\ref{eqn:Vphi_series}) as $V_0=A(1-\sqrt{B})H_0^2M_{\rm pl}^2$.
16
7
1607.01720
We study the nature of potentials in scalar field based models for dark energy—with both canonical and noncanonical kinetic terms. We calculate numerically, and using an analytic approximation around a ≈1 , potentials for models with constant equation-of-state parameter w<SUB>ϕ</SUB> . We find that for a wide range of models with canonical and noncanonical kinetic terms there is a simple approximation for the potential that holds when the scale factor is in the range 0.6 ≲a ≲1.4 . We discuss how this form of the potential can also be used to represent models with nonconstant w<SUB>ϕ</SUB> and, hence, how it could be used in reconstruction from cosmological data.
false
[ "cosmological data", "scalar field based models", "models", "potentials", "canonical and noncanonical kinetic terms", "dark energy", "reconstruction", "nonconstant w<SUB>ϕ</SUB", "both canonical and noncanonical kinetic terms", "a simple approximation", "state", "an analytic approximation", "a wide range", "the scale factor", "the potential", "the range", "a ≈1", "the nature", "this form", "We" ]
11.131959
0.426444
89
12513428
[ "Milgrom, Mordehai" ]
2016PhRvL.117n1101M
[ "Universal Modified Newtonian Dynamics Relation between the Baryonic and \"Dynamical\" Central Surface Densities of Disc Galaxies" ]
42
[ "Department of Particle Physics and Astrophysics, Weizmann Institute, Rehovot 7610001, Israel" ]
[ "2016ApJ...832L...8M", "2016MNRAS.463.3637R", "2016arXiv161007538M", "2016arXiv161009181K", "2017A&A...603A..11P", "2017A&A...603A..65T", "2017Galax...5...17D", "2017IJMPD..2650118C", "2017MNRAS.472..765T", "2017PhRvD..95f3020C", "2017RAA....17...74C", "2017arXiv170306110M", "2018ApJ...853...60W", "2018JCAP...03..038F", "2018PhRvD..97b3027C", "2018PhRvD..98f4015D", "2018arXiv180405840M", "2018arXiv180810545B", "2019MNRAS.483L..64D", "2019MNRAS.484.1421C", "2019MNRAS.487.2148G", "2020MNRAS.492.5865C", "2020PDU....2800468Z", "2020PDU....2800478C", "2020SHPMP..71..170M", "2021ApJ...923...68O", "2021CQGra..38n5022C", "2021PhRvD.103b3523B", "2021PhRvD.103d4043M", "2021SHPSA..88..193M", "2021SHPSA..88..220M", "2022NatAs...6...35L", "2022PASJ...74.1441C", "2022PDU....3801142C", "2022PhRvD.105h3003C", "2023InJPh..98.1923M", "2023MNRAS.518.6238C", "2023MNRAS.519.1277C", "2024InJPh..98.1923M", "2024JCAP...04..020M", "2024JCAP...05..042D", "2024arXiv240511576M" ]
[ "astronomy", "physics" ]
11
[ "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology" ]
[ "1963ApJ...138..385T", "1983ApJ...270..365M", "1983ApJ...270..371M", "1984ApJ...286....7B", "1995MNRAS.276..453B", "2002JPhA...35.1437M", "2009MNRAS.397.1169D", "2009MNRAS.398.1023M", "2010MNRAS.403..886M", "2012LRR....15...10F", "2014MNRAS.437.2531M", "2014SchpJ...931410M", "2015MNRAS.452.3650O", "2016ApJ...827L..19L" ]
[ "10.1103/PhysRevLett.117.141101", "10.48550/arXiv.1607.05103" ]
1607
1607.05103_arXiv.txt
Introduction} MOND \cite{milgrom83,milgrom83a} attributes the mass discrepancies in galactic systems not to dark matter but to a departure from the standard dynamics at low accelerations. Its basic tenets are: (i) Galactic systems showing large mass discrepancies are governed by new dynamics that are invariant under space-time scaling $(\vr,t)\rar\l(\vr,t)$. There, Newton's $G$ is replaced by a scale-invariant constant $\azg$. (ii) The boundary between the standard and the scale-invariant regime is marked by the constant $\az=\azg/G$, which is an acceleration. Thus, much below $\az$ -- `the deep-MOND limit' -- dynamics are scale invariant, and much above it standard dynamics are approached. References \cite{fm12,milgrom14c} are recent reviews of MOND. \par In contradistinction from the dark-matter paradigm, MOND decrees, as a law of physics, that baryons determine the full dynamics of a system. So, MOND predicts that various `dynamical' galaxy properties -- deduced from the measured accelerations -- are tightly correlated with `baryonic' properties -- those deduced directly from the distribution of baryonic mass (see, e.g., the recent Ref. \cite{milgrom14}). An example of such a `MOND law' is the mass-asymptotic-speed relation (MASSR) \cite{milgrom83a}, which is arguably the most famous prediction of an exact, functional relation. It had predicted a specific version of the `baryonic Tully-Fisher relation'. \par Here, I deal with a new MOND law: {\it a functional relation} between a `baryonic' and a `dynamical' property of pure disc galaxies. The one is the central surface (baryonic) density of the disc, $\Sbz$, the other is the total, `dynamically' measured central surface density, $\Sdz$. The impetus to look for such a MOND central-surface-densities relation (CSDR) has come from the recent finding of a correlation between two similar quantities in a large sample of disc galaxies in Ref. \cite{lelli16}. \par This MOND CSDR is quite different from other MOND laws of galactic dynamics -- such as the MASSR, or the discrepancy-acceleration relation -- and thus broadens the scope of the MOND codex (see Sec. \ref{discussion}). \par Preliminary aspects of the MOND CSDR have been discussed in Ref. \cite{milgrom09} in an approximate way, in terms of the mean density and characteristic size of the galaxy. In particular, it was shown that MOND predicts that, for $\Sbz\gg\SM\equiv\az/\tpg$, we have $\Sdz-\Sbz\approx\SM$, which implies that $\Sdz/\Sbz\approx 1$. For $\Sbz\ll\SM$, MOND predicts that $\Sdz$ scales as $(\SM\Sbz)^{1/2}$, with the coefficient estimated approximately, and being somewhat system dependent.\footnote{This prediction was made contrary to claims, e.g. in Ref. \cite{donato09}, tantamount to $\Sdz$ being independent of $\Sbz$ in this limit.} These predictions are born out by the findings of Ref. \cite{lelli16}. \par Here, I go rather far beyond these preliminary predictions, and show in Sec. \ref{derivation} that for pure discs, the MOND CSDR, $\Sbz$-$\Sdz$ relation is functional, and I derive $\Sdz(\Sbz)$ analytically for the full $\Sbz$ range. The predicted relation is compared with the data of Ref. \cite{lelli16} in Sec. \ref{data}. Section \ref{discussion} is a discussion.
Discussion} The MOND CSDR differs from other MOND laws (discussed, e.g., in Ref. \cite{milgrom14}) in several important regards: a. It is a relation between a `local' baryonic attribute, $\Sbz$, defined at the center, and a `global' dynamical one: the column dynamical density along the symmetry axis. b. It encompasses the full gamut of accelerations. In comparison, the MASSR relates a global baryonic attribute -- the total mass -- with the asymptotic rotational speeds, and it involves only the deep-MOND regime, and only the outskirts of systems. The MOND `mass-discrepancy-acceleration relation' for disc galaxies, relates the baryonic and dynamical accelerations at the same radius, in the plane of disc galaxies. And, the MOND mass-velocity-dispersion relation (an analog of the virial theorem) relates two global attributes, and holds in the low-acceleration regime. \par With good enough data, one can, in principle determine the MOND interpolating function by differentiating the observed $\SS(y)$, from eq. (\ref{vi}). This can also be done in an independent way from the discrepancy-acceleration relation. But this will have to await better data to constrain $\SS(y)$. \par The MOND CSDR falls very near the 3-parameter best-fit of Ref. \cite{lelli16}. The scatter around this best fit was shown in Ref. \cite{lelli16} to be $\approx 0.2$ dex, consistent with observational and procedural errors, i.e., with no intrinsic scatter. Furthermore, the residuals do not correlate with various galaxy properties, such as size, gas fraction, or stellar mass. \par This is exactly what MOND predicts, while, as also stressed in Ref. \cite{lelli16}, in the dark-matter paradigm, there is no reason why $\Sdz$, which would stand for the central column density of baryon plus dark matter, should be so well correlated with the local $\Sbz$. Especially, the correlation in the low $\Sbz$ region -- where the discrepancy between $\Sdz$ and $\Sbz$ is large -- goes quite against the grain of the cold-dark-matter paradigm: Even with schemes, such as `feedback', `abundance matching', and the like -- put in by hand to save this paradigm from various embarrassments -- one would expect large scatter in any relation between the `dynamical' and baryonic properties, which, to boot, one would expect to depend on galaxy properties (see, e.g., the relevant discussion in Ref. \cite{oman15}).
16
7
1607.05103
I derive a new modified Newtonian dynamics (MOND) relation for pure-disc galaxies: The "dynamical" central surface density, Σ<SUB>D</SUB><SUP>0</SUP> , deduced from the measured velocities, is a universal function of only the true, "baryonic'' central surface density, Σ<SUB>B</SUB><SUP>0</SUP> : Σ<SUB>D</SUB><SUP>0</SUP>=Σ<SUB>M</SUB>S (Σ<SUB>B</SUB><SUP>0</SUP>/Σ<SUB>M</SUB>) , where Σ<SUB>M</SUB>≡a<SUB>0</SUB>/2 π G is the MOND surface density constant. This surprising result is shown to hold in both existing, nonrelativistic MOND theories. S (y ) is derived: S (y )=∫<SUB>0</SUB><SUP>y</SUP>ν (y<SUP>'</SUP>)d y<SUP>'</SUP> , with ν (y ) being the interpolating function of the theory. The relation aymptotes to Σ<SUB>D</SUB><SUP>0</SUP>=Σ<SUB>B</SUB><SUP>0</SUP> for Σ<SUB>B</SUB><SUP>0</SUP>≫Σ<SUB>M</SUB> , and to Σ<SUB>D</SUB><SUP>0</SUP>=(4 Σ<SUB>M</SUB>Σ<SUB>B</SUB><SUP>0</SUP>)<SUP>1 /2</SUP> for Σ<SUB>B</SUB><SUP>0</SUP>≪Σ<SUB>M</SUB> . This study was prompted by the recent finding of a correlation between related attributes of disc galaxies by Lelli et al.. The MOND central-surface-densities relation agrees very well with these results.
false
[ "SUB", "Σ<SUB", "Σ<SUB>D</SUB><SUP>0</SUP>=(4 Σ<SUB>M</SUB>Σ<SUB", "Lelli et al", "Σ<SUB>B</SUB><SUP>0</SUP>≫Σ<SUB", "Σ<SUB>B</SUB><SUP>0</SUP>≪Σ<SUB", "disc galaxies", "Σ<SUB>D</SUB><SUP>0</SUP>=Σ<SUB>B</SUB><SUP>0</SUP", "y", "Σ<SUB>D</SUB><SUP>0</SUP>=Σ<SUB>M</SUB>S", "related attributes", "The MOND central-surface-densities relation", "ν", "S", "pure-disc galaxies", "both existing, nonrelativistic MOND theories", "G", "D</SUB><SUP>0</SUP", "Lelli", "The \"dynamical\" central surface density" ]
10.090203
2.106674
51
14285808
[ "Simpson, Fergus", "Jimenez, Raul", "Pena-Garay, Carlos", "Verde, Licia" ]
2018PDU....20...72S
[ "Dark energy from the motions of neutrinos" ]
11
[ "ICC, University of Barcelona (UB-IEEC), Marti i Franques 1, 08028, Barcelona, Spain", "ICC, University of Barcelona (UB-IEEC), Marti i Franques 1, 08028, Barcelona, Spain", "Instituto de Fisica Corpuscular, CSIC-UVEG, P.O. 22085, Valencia, 46071, Spain", "ICC, University of Barcelona (UB-IEEC), Marti i Franques 1, 08028, Barcelona, Spain" ]
[ "2016PhRvD..94j3518C", "2016arXiv160800493M", "2016arXiv161105710S", "2020IJMPD..2950106R", "2020PhRvD.101l3505K", "2020arXiv201008181R", "2021PDU....3100777K", "2021PDU....3400897K", "2021arXiv210507973R", "2022JHEAp..34...49A", "2024PhLB..85338687C" ]
[ "astronomy", "physics" ]
3
[ "Neutrinos", "Dark energy", "Interactions in the dark sector", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]
[ "1959PhRv..115..485A", "1981PhLB...98..265C", "1981PhLB...99..411G", "1982PhRvL..48.1522F", "1983NuPhB.219...31G", "1983RSPTA.310..347C", "1987PhRvL..59.2607W", "1988PhRvD..37.3406R", "1990PhRvD..42..293C", "1992PhR...215..203M", "1994PhRvD..49.5925B", "1995MNRAS.274L..73E", "1995NuPhB.449...25B", "1999PhRvD..59l3504S", "1999PhRvL..82..896Z", "2000PhRvD..62b3004K", "2000PhRvD..62d3511A", "2001NuPhS..91....3W", "2003PhRvD..67g3015F", "2004PhRvL..93i1801K", "2005PhRvD..71j3523T", "2005PhRvD..72f5024A", "2006PhRvD..74d3517M", "2006PhRvL..96f1301B", "2007MNRAS.379.1067P", "2007PhRvD..75b3502S", "2008JCAP...01..026E", "2008JCAP...07..020V", "2008PhRvD..78b3015A", "2008PhRvD..78b3505B", "2009JCAP...07..034G", "2009PhRvD..79a5018B", "2009PhRvD..79d3522C", "2009PhRvD..79d3526J", "2010MNRAS.402.2355V", "2010PhRvD..82h3505S", "2010PhRvD..82l3001P", "2011MNRAS.418..214B", "2011arXiv1109.5515G", "2012JInst...7.6001A", "2012PhRvC..86b1601G", "2013PhRvD..88h3505P", "2013PhRvL.110q4301G", "2013arXiv1308.1633M", "2013arXiv1310.4692C", "2014JCAP...07..046A", "2014PhRvD..90l3533C", "2015JCAP...07..014F", "2015PhRvD..91h3537S", "2015PhRvD..92b5037B", "2016PhRvD..94d3518P", "2016PhRvL.116h1301P" ]
[ "10.1016/j.dark.2018.04.002", "10.48550/arXiv.1607.02515" ]
1607
1607.02515_arXiv.txt
\begin{table*} \begin{tabular}{ |l|c|l| } \hline \bf{Event} & \bf{Scale Factor} & \bf{ Correlations} \\ \hline 1) Matter-Radiation Equality & $10^{-3.5}$ & \\ 2) Recombination & $10^{-3}$ & \\ 3) Neutrinos become non-relativistic & $10^{-2}$ & \\ 4) Dark Energy - Matter Equality & $1$ & \\ \hline 5) Neutrino-Radiation Equality & $10^{-2}$ & Caused by (3) \\ 6) Nonlinear Density Perturbations & $10^{-1}$ & Requires (1) \\ 7) Our Existence & 1 & Requires (6) \\ \hline \end{tabular} \caption{A selection of recent cosmic events. A minimum of three distinct events is required in order for the four fluids to reverse their positions in the density hierarchy. These are represented here by three moments of equality. Known correlations limit the number of independent events in this table to four. As-yet undiscovered correlations may further reduce this number.} \label{tab:timeline} \end{table*} In the standard Lambda Cold Dark Matter (\lcdm) paradigm, each of the four key components (photons, neutrinos, dark matter, dark energy) crossed paths within the past ten e-folds. This corresponds to seven points of equality, in addition to other notable events such as recombination and the onset of nonlinear structure formation. The traditional coincidence problem in cosmology, casually stated as ``why now?", relates to the close proximity in scale factor between our existence and the onset of dark energy domination. This is arguably the most straightforward coincidence to resolve - we are better described as a sample in time rather than a sample in scale factor. Therefore if the Universe were to recollapse or rip within the next trillion years, our existence is not especially close to the onset of dark energy. Some of the other coincidences appear more stubborn, however. Why did the four major contributors to the cosmic energy budget all cross paths with each other in rapid succession? The seemingly congested cosmic timeline, as summarised in Table \ref{tab:timeline}, can be remedied if some events are correlated with each other. For example, we should not be surprised to find that nonlinear structure formation initiated within a few e-folds of our existence, since it is a prerequisite for complex life. Furthermore, nonlinear structure formation was catalysed by the onset of matter domination, so these should not be interpreted as wholly independent events. Can we reduce this list of four independent events still further? In particular, could the onset of dark energy be a direct consequence of another recent event? Upper bounds on the coupling strength between dark energy, dark matter, and neutrinos are extremely weak \citep{simpscat}. It is therefore feasible that non-gravitational interactions were responsible for the onset of cosmic acceleration. A number of models have been proposed which invoke energy exchange between a scalar field and a second fluid \citep{2000PhRvD..62d3511A}. Motivated by the approximate equivalence of the neutrino mass and the dark energy density, a direct connection between the two has already been investigated in some detail \citep{fardon2004dark, kaplan2004neutrino, 2006Brookfield}. However these models tend to suffer from instabilities in their perturbations \cite{2005Afshordi}. Recently a class of models involving a derivative coupling at the fluid level were presented in Ref.~\cite{2013Pourtsidou}, and explored further by Refs.~\cite{2015Skordis, 2016Pourtsidou, 2016BaldiSimpson}. In this work we shall first explore the means by which energy exchange can facilitate the formation of dark energy. Then we present the particular case of a derivative coupling at the particle level, between neutrinos and a scalar field. \begin{figure}[b] \includegraphics[width=80mm]{fig1} \caption{ The energy densities associated with the four major cosmological fluids (photons, neutrinos, dark matter, and dark energy) in units of the present day critical density. The solid line represents a toy model which demonstrates how the dark energy could acquire its present day value, while having a negligible impact on the other fluids. \label{fig:densities}} \end{figure}
Conventional models of quintessence generally suffer from the same kind of fine-tuning problems as the cosmological constant they were designed to usurp. Ordinarily, the scalar must either lie at a local minimum of a finely-tuned potential, or be at a finely-tuned location within an extremely flat potential. However we have demonstrated that if a scalar field is in the process of energy exchange, it can mimic a cosmological constant even when displaced from the local minimum of its potential. This allows potentials of \emph{any} shape to act like a cosmological constant. The underlying phenomenology of this freezing process is founded in classical mechanics. A cyclist who is pedalling uphill might begin to struggle against an increasingly steep gradient. Yet no matter how feeble the rider's power output becomes, they can always maintain a constant pedalling rate, simply by selecting an appropriately high gear. They can \emph{never} roll backwards. We have demonstrated that a derivative coupling between neutrinos and the majoron could provide the required energy transfer. As neutrinos become non-relativistic, their velocities begin to align with the gradients in the local gravitational potential. In the presence of a derivative coupling, this creates an effective potential for the scalar field. The effective potential can act as either the driving force or a brake, depending on the sign of the coupling constant. In each case, once the freezing process sets in, the field's equation of state rapidly approaches that of a cosmological constant. One path towards resolving these issues may be anthropic selection from within an ensemble: a multiverse. This has been touted as an explanation for the value of the cosmological constant \citep{1987PhRvL..59.2607W, 1995MNRAS.274L..73E, 2007PeacockAnthropic, 2015Piran} and the neutrino mass \cite{2005Tegmark, 2015Bousso}. However, before accepting a hypothesis which is not experimentally falsifiable, it is wise to exclude all viable alternatives which \emph{are}. It has been argued in the past \cite{2001Witten} that any theory of quantum gravity will not allow B-L to be a symmetry of nature and thus could provide neutrinos with their mass. It is, therefore, curious that dark energy could indirectly result from broken symmetries, due to the quantum nature of gravity; the so-called IR-UV connection. In summary, the freezing mechanism presented in this work offers a generic means to generate dark energy from any scalar potential $V(\phi)$. The derivative coupling model we propose offers a natural and experimentally verifiable solution to the dark energy coincidence problem.
16
7
1607.02515
Ordinarily, a scalar field may only play the role of dark energy if it possesses a potential that is either extraordinarily flat or extremely fine-tuned. Here we demonstrate that these restrictions are lifted when the scalar field undergoes persistent energy exchange with another fluid. In this scenario, the field is prevented from reversing its direction of motion, and instead may come to rest while displaced from the local minimum of its potential. Therefore almost any scalar potential is capable of initiating a prolonged phase of cosmic acceleration. If the rate of energy transfer is modulated via a derivative coupling, the field undergoes a rapid process of freezing, after which the field's equation of state mimicks that of a cosmological constant. We present a physically motivated realisation in the form of a neutrino-majoron coupling, which avoids the dynamical instabilities associated with mass-varying neutrino models. Finally we discuss possible means by which this model could be experimentally verified.
false
[ "cosmic acceleration", "dark energy", "mass-varying neutrino models", "state mimicks", "energy transfer", "freezing", "rest", "a cosmological constant", "motion", "possible means", "a scalar field", "the scalar field", "a neutrino-majoron coupling", "a derivative coupling", "a prolonged phase", "the local minimum", "a potential", "its potential", "the dynamical instabilities", "the field" ]
7.340179
-1.160099
-1
12593474
[ "Rau, Markus Michael", "Hoyle, Ben", "Paech, Kerstin", "Seitz, Stella" ]
2017MNRAS.466.2927R
[ "Correcting cosmological parameter biases for all redshift surveys induced by estimating and reweighting redshift distributions" ]
7
[ "Ludwig-Maximilians-Universität München, Universitäts-Sternwarte, Scheinerstr. 1, D-81679 Munich, Germany; Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany", "Ludwig-Maximilians-Universität München, Universitäts-Sternwarte, Scheinerstr. 1, D-81679 Munich, Germany", "Ludwig-Maximilians-Universität München, Universitäts-Sternwarte, Scheinerstr. 1, D-81679 Munich, Germany; Excellence Cluster Universe, Bolzmannstr. 2, D-85748 Garching, Germany", "Ludwig-Maximilians-Universität München, Universitäts-Sternwarte, Scheinerstr. 1, D-81679 Munich, Germany; Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany" ]
[ "2019MNRAS.485.3642H", "2019MNRAS.486.2730A", "2020MNRAS.491.4768R", "2020MNRAS.496.4769H", "2021A&A...655A..44E", "2022MNRAS.509.4886R", "2023MNRAS.524.5109R" ]
[ "astronomy" ]
5
[ "catalogues", "surveys", "galaxies: distances and redshifts", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1985AJ.....90..418K", "1992mde..book.....S", "1996ApJ...461L..65F", "1998PhRvL..81.2004K", "2000ApJ...536..571B", "2002PhRvD..65f3001H", "2005IJMPA..20.3121F", "2008ApJ...684...88N", "2009A&A...504..359J", "2009A&A...507..105J", "2009MNRAS.396.2379C", "2010ApJ...720.1351H", "2013ApJ...764..167S", "2013MNRAS.431.1547B", "2013MNRAS.432.1483C", "2013MNRAS.432.2945H", "2013Msngr.154...44D", "2014MNRAS.444..129C", "2014MNRAS.445.1482S", "2015A&C....12...45Z", "2015MNRAS.447.2961S", "2015MNRAS.449.1043B", "2015MNRAS.450..305H", "2015MNRAS.452.3710R", "2016MNRAS.459..971K", "2016MNRAS.460..163R", "2016MNRAS.460.1371B", "2016MNRAS.460.4258L", "2016MNRAS.461.3432S", "2016MNRAS.463.3737C", "2016PhRvD..94d2005B" ]
[ "10.1093/mnras/stw3338", "10.48550/arXiv.1607.00383" ]
1607
1607.00383_arXiv.txt
16
7
1607.00383
Photometric redshift uncertainties are a major source of systematic error for ongoing and future photometric surveys. We study different sources of redshift error caused by choosing a suboptimal redshift histogram bin width and propose methods to resolve them. The selection of a too large bin width is shown to oversmooth small-scale structure of the radial distribution of galaxies. This systematic error can significantly shift cosmological parameter constraints by up to 6σ for the dark energy equation-of-state parameter w. Careful selection of bin width can reduce this systematic by a factor of up to 6 as compared with commonly used current binning approaches. We further discuss a generalized resampling method that can correct systematic and statistical errors in cosmological parameter constraints caused by uncertainties in the redshift distribution. This can be achieved without any prior assumptions about the shape of the distribution or the form of the redshift error. Our methodology allows photometric surveys to obtain unbiased cosmological parameter constraints using a minimum number of spectroscopic calibration data. For a DES-like galaxy clustering forecast, we obtain unbiased results with respect to errors caused by suboptimal histogram bin width selection, using only 5k representative spectroscopic calibration objects per tomographic redshift bin.
false
[ "redshift error", "tomographic redshift bin", "unbiased cosmological parameter constraints", "cosmological parameter constraints", "suboptimal histogram bin width selection", "Photometric redshift uncertainties", "systematic error", "bin width", "errors", "bin", "spectroscopic calibration data", "a suboptimal redshift histogram bin width", "unbiased results", "the redshift error", "commonly used current binning approaches", "systematic and statistical errors", "the redshift distribution", "methods", "galaxies", "photometric surveys" ]
12.115962
4.014999
157
14285822
[ "Kosovichev, Alexander G.", "Zhao, Junwei", "Ilonidis, Stathis" ]
2018vsss.book...15K
[ "Local Helioseismology of Emerging Active Regions: A Case Study" ]
3
[ "New Jersey Institute of Technology", "Stanford University", "Stanford University" ]
[ "2017ApJ...839...63K", "2019LRSP...16....3T", "2021arXiv210901833R" ]
[ "astronomy", "physics" ]
3
[ "Sun: activity", "Sun: helioseismology", "Sun: interior", "Sun: magnetic fields", "Sun: oscillations", "Astrophysics - Solar and Stellar Astrophysics", "Physics - Fluid Dynamics", "Physics - Plasma Physics" ]
[ "1979ApJ...230..905P", "1996ApJ...461L..55K", "1997ASSL..225..241K", "1999PCEA...24..215J", "2000SoPh..192..193B", "2001ApJ...557..384Z", "2001ApJ...561L.229B", "2001ESASP.464..187B", "2002ApJ...570..855H", "2002ApJ...571..966G", "2003ApJ...591..446Z", "2003ESASP.517..103H", "2004ApJ...603..776Z", "2004ApJ...608..580B", "2004ApJ...614.1073B", "2006ApJ...640..516C", "2006SSRv..124....1K", "2007AN....328..228B", "2009ASPC..415..411Z", "2009ApJ...698.1749H", "2010ApJ...708..304Z", "2010ApJ...719..307K", "2012SoPh..275..207S", "2012SoPh..275..357C", "2012SoPh..275..375Z", "2013ApJ...762..131B", "2013ApJ...774L..29Z", "2013ApJ...777..138I", "2014ApJ...789...35F", "2014ApJ...789L...7Z", "2016A&A...588A.150K", "2016ApJ...819..104G", "2016ApJ...822L..22M", "2016LNP...914...25K" ]
[ "10.48550/arXiv.1607.04987" ]
1607
1607.04987_arXiv.txt
\label{sec1} Emergence and formation of magnetic active regions on the surface of the Sun is one of the central problems of solar physics. It is of the fundamental importance in astrophysics because active regions are one of the primary manifestations of the solar and stellar magnetism. In addition, solar active regions (AR) are the major drivers of the solar variability, geospace and planetary space environments and space weather. Understanding of the emergence and evolution of active regions is a key to developing the knowledge and capability to detect and predict extreme conditions in space. The uninterrupted helioseismology and magnetic data from Solar Dynamics Observatory provide unique opportunities for comprehensive studies that can uncover the basic mechanisms of active region formation, evolution, and their flaring and CME activity \citep{Scherrer2012}. Recent studies revealed that the emergence and magnetic structure of active regions are closely linked to the plasma flows on the surface and in the subsurface layers \citep{Hindman2009,Komm2011,Komm2012,Kosovichev2012,Kosovichev2006,Birch2013}. For example, diverging subsurface flows have been detected prior the emergence of active regions, shearing and twisting flows are found to be associated with flaring activity, large-scale converging flows formed around active regions affect the meridional circulation and along with the tilt of active regions (Joy’s law) are believed to be among primary factors determining the strength and duration of the future solar cycles. The conventional wisdom is that the active regions are a result of emergence of toroidal magnetic flux ropes formed near the bottom of the convection zone by a dynamo process. This theory can explain the Joy's law and the absence of active regions at high latitudes providing the initial magnetic field strength is about 60~kG \citep{DSilva1993}, which greatly exceeds the equipartition field strength and has not been reproduced by the dynamo theories. Current 3D MHD global-Sun models computed in the anelastic approximation have shown that the dynamo-generated field can be organized in the form of flux tubes but on much larger scale \citep{Brun2004,Fan2014,Guerrero2016}. These models indicates that the active regions and sunspots are probably formed in the near-surface layers, but the anelastic approximation becomes invalid close to the surface, where compressibility effects play significant role. With the currently available computational resources the realistic compressible radiative MHD simulations are capable to model only relatively shallow near-surface. These simulations have revealed a process of spontaneous formation of compact pore-like structures from initially distributed magnetic fields, maintained by converging downdrafts, however, other mechanisms of magnetic self-organization may be also involved \citep{Kapyla2016,Kitiashvili2010,Masada2016}. The magnetic self-organization process is probably a key to understand the formation of sunspots and active regions. It involves a complex interaction of turbulence with magnetic field, but in all cases the large-scale flow pattern includes compact regions of converging downdrafts around magnetic structures in shallow $\sim 5$~Mm deep regions. In the deeper layers the flows are mostly diverging. This flow pattern corresponds to the Parker's cluster model of sunspots. It has been observed by the time-distance helioseismology analysis of the SOHO/MDI and Hinode/SOT data \citep{Zhao2001,Zhao2009b,Zhao2003}. The data analysis also showed that in the decaying sunspots the flows become diverging. Other helioseismology methods, such as the ring-diagram analysis and the helioseismic holography, provide the subsurface flow maps with a lower resolution than the time-distance helioseismology, and did not confirm the existence of the converging downdrafts beneath the sunspots. Instead, they inferred diverging flows of a larger scale around sunspots over the whole depth range probed by these techniques \citep{Hindman2009}. The current helioseismology measurements in regions of strong magnetic field are subject to significant systematic errors due to the uncertainties in the Doppler-shift measurements, large variations of the sound speed causing non-linear wave effects, non-uniform distribution of acoustic sources and MHD wave transformation effects. These uncertainties mostly affect inferences of the sound-speed distribution beneath the sunspots \citep[for a recent review see][]{Kosovichev2012}. The helioseismic inferences of subsurface flows are based on measuring and inverting the travel-time differences for the waves traveling along the same path in the opposite directions. Such reciprocal signals are less sensitive to the systematic uncertainties. However, a complete testing and calibration of the flow inferences based on numerical simulations of wave propagation in sunspots has not been completed. In this work we mostly focus on flows of active regions that are formed beneath sunspots, and, in particular, on the flow patterns during the formation and evolution of a large emerging active region. The primary goal is to investigate the process of formation of the large-scale converging flows that affect the meridional circulation and magnetic flux transport. As a case study we consider a large emerging active region NOAA 11726.
As a case study we presented a local helioseismology analysis of the subsurface dynamics of emerging active region NOAA~11726 which is the largest emerging region observed by the SDO/HMI instrument during the first five years of operation. The active region emergence was detected at the depth of $62-75$~Mm about 12 hours before the first bipolar magnetic structure appeared on the surface, and 2 days before the emergence of most of the magnetic flux. The characteristic speed of emergence estimated from the signal delay in two layers, $62-75$~Mm and $42-55$~Mm deep, is about 1.4~km/s. During emergence of the initial bipolar structure no specific large-scale flow pattern in the depth range $0-20$~Mm were identified. Nevertheless, the region of the flux emergence is characterized by an enhanced horizontal flow divergence that corresponds to the spatial separation of the magnetic polarities. The speed of emergence determined by tracking the initial divergence signal with depth is about $1.4$~km/s, very close to the emergence speed in the deep layers. As the emerging magnetic flux becomes concentrated in sunspots local converging flows are observed beneath the forming sunspots. The converging flows are most prominent in the depth range $1-3$~Mm, and remain converging after the formation process is completed. The structure of the converging flows is complicated and apparently reflects the sunspot structural evolution and interaction with the surrounding convection flows. The characteristic speed of these flows is about 0.3~km/s. In the deeper layers the flows beneath the sunspots are predominantly diverging and occupy larger areas. At the depth about 15~Mm the diverging flows occupy a large area around the whole active region. By applying a Gaussian filter to smooth the supergranulation-scale flows we investigated the formation of large-scale converging flows around the active region. The scale of these flows is much larger than the size of the active region, and the typical flow speed is about 30~m/s. The formation of the converging flows appears as a diversion of the zonal shearing flows towards the active region, accompanied by formation of a large-scale vortex structure. This process occurs when a substantial amount of the magnetic flux emerged on the surface, and the converging flow pattern remains stable during the following evolution of the active region. The flow helicity is opposity in sign to the subsurface hemispheric helicity mostly determined by the supergranulation flows, but does not significantly contributes to the helicity balance. The primary effect of the converging large-scale flows is in changing the speed of the mean meridional flow in the top 10~Mm-deep layer. The synoptic flow maps presented for the Carrington rotation that includes our case study of AR~11726 show that the large-scale subsurface inflows are typical for other active regions. In the deeper layers ($10-13$~Mm) the flows become diverging, and quite strong beneath some active regions. The active region 11726 in the synoptic map is presented at the beginning of the emergence, but in the deep layers is accompanied by an area of strong divergence, which is off-side of the emerged flux. It remains to be seen if the deep diverging flows indicate on the future development of active regions, but the synoptic map of the following rotation shows that the active region continued to grow on the far side of the Sun and became very large (it received the new NOAA number 11745). In addition to the flows around active regions the synoptic maps reveal a complex evolving pattern of large-scale flows on the scale much larger than supergranulation. It appears that these flows correlate with the large-scale magnetic field outside active region. The exact relationship has not been established, but the presented case study encourages further in-depth investigations of the solar subsurface dynamics, both observationally and by numerical MHD simulations.
16
7
1607.04987
Local helioseismology provides a unique opportunity to investigate the subsurface structure and dynamics of active regions and their effect on the large-scale flows and global circulation of the Sun. We use measurements of plasma flows in the upper convection zone, provided by the Time-Distance Helioseismology Pipeline developed for analysis of solar oscillation data obtained by Helioseismic and Magnetic Imager (HMI) on Solar Dynamics Observatory (SDO), to investigate the subsurface dynamics of emerging active region NOAA 11726. The active region emergence was detected in deep layers of the convection zone about 12 hours before the first bipolar magnetic structure appeared on the surface, and 2 days before the emergence of most of the magnetic flux. The speed of emergence determined by tracking the flow divergence with depth is about 1.4 km/s, very close to the emergence speed in the deep layers. As the emerging magnetic flux becomes concentrated in sunspots local converging flows are observed beneath the forming sunspots. These flows are most prominent in the depth range 1-3 Mm, and remain converging after the formation process is completed. On the larger scale converging flows around active region appear as a diversion of the zonal shearing flows towards the active region, accompanied by formation of a large-scale vortex structure. This process occurs when a substantial amount of the magnetic flux emerged on the surface, and the converging flow pattern remains stable during the following evolution of the active region. The Carrington synoptic flow maps show that the large-scale subsurface inflows are typical for active regions. In the deeper layers (10-13 Mm) the flows become diverging, and surprisingly strong beneath some active regions. In addition, the synoptic maps reveal a complex evolving pattern of large-scale flows on the scale much larger than supergranulation
false
[ "active region", "active regions", "emerging active region", "sunspots local converging flows", "flows", "plasma flows", "The active region emergence", "large-scale flows", "some active regions", "the active region", "deep layers", "Solar Dynamics Observatory", "the larger scale converging", "emergence", "solar oscillation data", "the converging flow pattern", "Magnetic Imager", "dynamics", "the large-scale flows", "the first bipolar magnetic structure" ]
12.046253
15.671393
2
12516676
[ "Chelpanov, A. A.", "Kobanov, N. I.", "Kolobov, D. Y." ]
2016SoPh..291.3329C
[ "Influence of the Magnetic Field on Oscillation Spectra in Solar Faculae" ]
8
[ "Institute of Solar-Terrestrial Physics of Siberian Branch of Russian Academy of Sciences, Irkutsk, Russia", "Institute of Solar-Terrestrial Physics of Siberian Branch of Russian Academy of Sciences, Irkutsk, Russia", "Institute of Solar-Terrestrial Physics of Siberian Branch of Russian Academy of Sciences, Irkutsk, Russia" ]
[ "2018SoPh..293..157C", "2019A&A...627A..10R", "2019Ap&SS.364...29S", "2020AstBu..75..176S", "2020Ge&Ae..59..904E", "2020SoPh..295...94C", "2021A&A...652A..43Y", "2022SoPh..297...52C" ]
[ "astronomy" ]
6
[ "Active regions", "magnetic fields", "Magnetic fields", "photosphere", "Oscillations", "solar", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1965ApJ...141.1131O", "1967SoPh....2....3H", "1969SoPh....9...35P", "1977A&A....55..239B", "1995itsa.conf..387M", "1997ApJ...474..810M", "1998BAMS...79...61T", "1999ASPC..184..136N", "2001A&A...379.1052K", "2002A&A...386.1123S", "2005A&A...437.1055M", "2006ApJ...647L..77M", "2006ApJ...648L.151J", "2006MNRAS.372..551S", "2007A&A...471..961M", "2007SoPh..246..273K", "2008A&A...481L..25I", "2008ApJ...676L..85K", "2009ApJ...702L.168D", "2010SoPh..262...35G", "2011A&A...525A..41K", "2011ARep...55..532K", "2011AstL...37..799T", "2011SoPh..268..329K", "2011SoPh..270..175A", "2012A&A...545A..22K", "2012ApJ...746..119R", "2012PhyU...55A...4S", "2012RSPTA.370.3193D", "2012SoPh..275..229S", "2012SoPh..279..295L", "2013A&A...554A.146K", "2013A&A...559A.107K", "2014ARep...58..272K", "2014SoPh..289.3483H", "2015SoPh..290..363K" ]
[ "10.1007/s11207-016-0954-6", "10.48550/arXiv.1607.01877" ]
1607
1607.01877_arXiv.txt
\label{S-Introduction} \par Studying waves in the solar atmosphere can help answer the question of energy transfer to the outer layers of the Sun. Also, the observed characteristics of waves can provide information on the physical conditions at different heights of the atmosphere \citep{moortel, stepanov}. Oscillations in faculae have been of special interest for 50 years \citep{orrall, Howard}, since faculae are the most abundant solar magnetic activity manifestation, and they occupy a considerable fraction of the solar surface. \par The characteristics of the magnetic field in active regions have been linked to the observed oscillation characteristics. For example, in sunspots, high-frequency oscillations tend to concentrate in the central part of a spot, where magnetic field lines are closest to vertical, and low-frequency oscillations are observed in the outer sunspot regions, where the magnetic field is inclined \citep{Jefferies, McIntosh, Kobanov11a, reznikova, kobanov13}. Oscillations in microwave data are observed above regions with strong magnetic field \citep{abramov}. \par Faculae as well are characterized by different strength and inclination of the magnetic field, though the distribution of these parameters are more chaotic than those in sunspots (\citealp{ishikawa}; \citealp{pillet}; \citealp{guo}). There are, however, patterns in the oscillation observation results: low-frequency oscillations increase at the facula boundaries, while three- and five-minute oscillations are observed within faculae \citep{kobanov11}. \citet{wijn} registered waves with three-minute period propagate in the central facular chromosphere, and five-minute waves propagate at the peripheral regions. The influence of the magnetic field on oscillation characteristics in faculae was studied by \citet{khomenko, kostik12, kostik13}. \citet{kostik13} noted increase in the observed period in the regions with high magnetic field strength. \citet{turova} reported a strong five-minute period in the chromosphere of a facula at the base of a coronal hole, and decrease in the oscillation period with increasing height. \par Magnetic field configuration is a probable cause for such a distribution of oscillations in faculae. However, the magnetic field azimuthal angle in faculae seems somewhat random \citep{pillet}, and the connection between the oscillations at different height levels is difficult to trace \citep{kobanov11b}. \par In this work we analyse characteristics of intensity, Doppler velocity and magnetic field oscillations in facular magnetic knots and compare them to those in facular patches with magnetic field of a moderate strength.
% \label{S-Conclusions} \par i) Magnetic field strength oscillation power spectra in facular magnetic knots have prominent peaks at a frequency of about 4.8\,mHz, while spectra of the moderate strength patches lack such peaks. ii) The intensity and Doppler velocity spectral composition of the photospheric Fe\,\textsc{i}\,6173\,\AA\ line is constant over the area of faculae. In magnetic knots, however, the oscillation power is reduced 2--3-fold in both intensity and velocity signals. iii) At the upper-photospheric level of the 1700\,\AA\ line, the dominant frequency of the oscillations in knots is increased by 0.5--1\,mHz compared to the moderate-field patches. iv) In the transition region---the He\,\textsc{ii}\,304\,\AA\ line---the significant peaks of the magnetic knot oscillation spectra sit between 3 and 6\,mHz, and those of the peripheral region spectra sit below 3\,mHz. \par We believe that such parameters of the oscillation distributions are caused by the magnetic field configuration in faculae. At the low atmospheric heights, the magnetic field lines are close to vertical over the whole area of a faculae, and the spectral composition at these heights does not change significantly in different parts of a facula. The field lines are inclined at the facular periphery in the upper photosphere, and more so in the transition region. This causes the decrease in the cut-off frequency and enables low-frequency waves to propagate upward. \hyphenation{Pro-jects} \small{\textbf{Acknowledgements}. This study was supported by Project No.\,16.3.2 of ISTP SB RAS, by the Russian Foundation for Basic Research under grants 16-32-00268 mol\_a and No.\,15-32-20504 mol\_a\_ved. We acknowledge the NASA/SDO science team for providing the data. We are grateful to the highly skilled anonymous reviewer, whose valuable remarks and suggestions helped us greatly improve the paper.
16
7
1607.01877
We study oscillation parameters in faculae above magnetic knots and in the areas adjacent to them. Using Solar Dynamics Observatory (SDO) data, we analyse oscillations in magnetic field strength, Doppler velocity, and intensity signals for the lower photosphere, and in intensity for the higher levels. We find peaks in the oscillation spectra for the magnetic field strength in magnetic knots at a frequency of about 4.8 mHz, while there are no such frequencies in the adjacent facular patches, which have a moderate field strength. In contrast, the spectra for the Doppler velocity oscillations are similar for these types of regions: in both cases, the significant peaks are in the 2.5 - 4.5 mHz range, although the oscillations in magnetic knots are 2 - 3 times weaker than those at the facular periphery. In the upper photosphere, the dominant frequencies in magnetic knots are 0.5 - 1 mHz higher than in the medium-field regions. The transition region oscillations above magnetic knots are mainly concentrated in the 3 - 6 mHz range, and those above moderate-field patches are concentrated below 3 mHz.
false
[ "magnetic field strength", "magnetic knots", "the magnetic field strength", "oscillation parameters", "oscillations", "a moderate field strength", "moderate-field patches", "Doppler velocity", "regions", "intensity", "the adjacent facular patches", "the Doppler velocity oscillations", "The transition region oscillations", "Doppler", "peaks", "the medium-field regions", "Solar Dynamics Observatory", "the oscillation spectra", "the facular periphery", "the higher levels" ]
12.059834
15.46882
2
12437382
[ "Martinez", "Janin-Potiron" ]
2016arXiv160708351M
[ "Laser-guide-stars used for cophasing broad capture ranges" ]
0
[ "-", "-" ]
null
[ "astronomy" ]
5
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "1985A&A...152L..29F", "1998ApOpt..37..140C", "1999A&A...344L..45P", "2000A&AS..144..533A", "2005JGRD..11023302S", "2006SPIE.6267E..2YP", "2008SPIE.7012E..3CB", "2011ApOpt..50.2708V", "2012ApPhL.101x1109H", "2012JASTP..74..181P" ]
[ "10.48550/arXiv.1607.08351" ]
1607
1607.08351_arXiv.txt
Adaptive optics (AO) enables obtaining diffraction-limited imaging provided that a guiding light-source is used as a reference for the wavefront correction. Bright natural stars are required for this, but the sky coverage is unfavorably low, in particular in the visible wavelength range. To overcome this limitation, artificial laser-guide-stars (LGSs) have been proposed in the 1980s \citep{LGS85} to create a source that lies above the turbulent layers of the atmosphere. LGS-based AO systems are now routinely used in all leading observatories (e.g., Keck, Gemini, ESO/VLT, and Lick), and multiple LGS AO systems are intensively employed, as demonstrated by the four LGSs that simultaneously shone above the Unit 4 telescope (UT4) at ESO's VLT in early May 2016. Different types of lasers are suitable for AO applications, but the most recent LGS systems make use of mesospheric scattering of a sodium laser. The principle of the sodium guide star is to tune the wavelength of the laser radiation to a resonance of sodium atoms at 589\,nm (D2 line). Sodium atoms, naturally present in the mesosphere at an altitude of 95\,km (on average), will absorb laser light and subsequently emit fluorescence at the same wavelength. The intensity of the created star depends on the amount of sodium atoms present at the altitude of the mesospheric sodium layer. In this context, various astronomical and atmospherical studies have been conducted over the past decades to understand the temporal and spatial characteristics of the atmosphere sodium layer \citep[e.g., ][]{GE98, P1999, A2000, SLANGER2005, PLANE2012}. The future of ground-based astronomy in the next decades is bound to the next generation of extremely large telescopes (ELTs). ELTs represent a major change in dimension, wavefront control strategies, and execution time. The wavefront control of ELTs includes three main systems: the aforementioned AO system that corrects for the atmospheric turbulence, active optics that correct for misalignments and deformations of the telescope adaptive mirror, and phasing optics that correct for the misalignment of individual segments of the primary mirror. Generally, cophasing is a three-step procedure to correct for the initial misalignment imposed by mechanical structure constraints to the final alignment: (1) a mechanical pre-phasing step using hand-held optical tools, (2) a coarse-phasing step using a dedicated phasing sensor, and (3) a fine-phasing step at high precision using a dedicated phasing sensor. Step (1) will be critical for ELTs because their primary mirrors require hundreds of segments, several of which will frequently need to be substituted for recoating (the reflectivity coating lifetime is generally limited to 18 months). In this context, mechanical integration and step (1) must reduce the piston error to the capture range of actual optical cophasing sensors. Enlarging this capture range is indispensable to relax mechanical integration requirements and to speed up integration and re-integration process before the observing run. Conventional coarse-phasing techniques use multiple wavelength or coherence methods \citep[e.g., ][]{CHANAN98, VIGAN11} to increase the capture range of the phasing sensor, which is limited by the so-called $\pi$-ambiguity otherwise. These methods successfully enlarge the initial capture range of a cophasing sensor to $\sim10\, \mu m$ $rms$. With the method we propose, segments of a telescope can reliably be phased from a $rms$ piston error of more than $1000\, \mu m$. Our method is called the doublet-wavelength coherence technic (DWCT) and does not require any hardware, except for the LGSs offered by the telescope. The DWCT is a new solution based on multiple LGSs lasing at different wavelengths through the analysis of the coherence signature in the resulting cophasing sensor signal. Because it exceeds the actual extended capture range of the cophasing sensor by a factor of 100, it might eliminate the need of the man-made mechanical pre-phasing step (1) along with inherent stability and compatibility independent of the phasing sensor type. In the following sections, the general background of phasing optics is introduced, and conventional coarse measurement techniques (multiwavelength and coherence methods) are presented for pedagogical reasons and because they are considered the basis for this work. Then the theory of the proposed DWCT and its advantages and limitations are discussed. \label{introduction}
\label{conclusion} The simplified description given in this paper is not thoroughly satisfactory because LGS systems are rather complex. Nonetheless, the preliminary analysis of the DWCT with LGSs suggests that the approach is worthy of interest. By using both the coherence and the double-line properties, the DWCT simultaneously grants access to steps (1) and (2) in one single interferogram and thus represents an enhancement of the coherence method. Advantages are not only the broad capture range, but also the inherent stability of the method, because the response coherence curve is Gaussian, and not periodic. The method is well suited to the re-integration of segments (or integration of spare segments) on a daily basis that return from re-coating. The DWCT improves the phasing optics capture paradigm to the millimetric domain and makes the man-made mechanical pre-phasing step redundant % This will accelerate the coarse-phasing step before the final fine-phasing step. If that lasing on multiple sodium D lines can be implemented, then the DWCT is a powerful tool that can also be used with phasing sensor designs. Applications to multi-aperture interferometric arrays still need to be explored.
16
7
1607.08351
Segmented primary mirrors are indispensable to master the steady increase in spatial resolution. Phasing optics systems must reduce segment misalignments to guarantee the high optical quality required for astronomical science programs. Modern telescopes routinely use adaptive optics systems to compensate for the atmosphere and use laser-guide-stars to create artificial stars as bright references in the field of observation. Because multiple laser-guide-star adaptive optics are being implemented in all major observatories, we propose to use man-made stars not only for adaptive optics, but for phasing optics. We propose a method called the doublet-wavelength coherence technique (DWCT), exploiting the D lines of sodium in the mesosphere using laser guide-stars. The signal coherence properties are then used. The DWCT capture range exceeds current abilities by a factor of 100. It represents a change in paradigm by improving the phasing optics capture range from micrometric to millimetric. It thereby potentially eliminates the need of a man-made mechanical pre-phasing step. Extremely large telescopes require hundreds of segments, several of which need to be substituted on a daily basis to be recoated. The DWCT relaxes mechanical integration requirements and speeds up integration and re-integration process.
false
[ "adaptive optics systems", "adaptive optics", "artificial stars", "optics systems", "astronomical science programs", "mechanical integration requirements", "optics", "spatial resolution", "observation", "integration and re-integration process", "bright references", "segment misalignments", "multiple laser-guide-star adaptive optics", "laser guide-stars", "capture range", "current abilities", "man-made stars", "segments", "millimetric", "the phasing optics" ]
14.255198
10.50998
10
12521176
[ "Vianello, G.", "Omodei, N.", "Chiang, J.", "Digel, S." ]
2017ApJ...841L..16V
[ "Searching for High-energy Gamma-ray Counterparts to Gravitational-wave Sources with Fermi-LAT: A Needle in a Haystack" ]
6
[ "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", "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" ]
[ "2017ApJ...835...82R", "2017ApJ...846L...5G", "2017PhLB..773..219D", "2017arXiv171005450F", "2018ApJ...861...85A", "2019ApJ...878...52A" ]
[ "astronomy" ]
5
[ "gamma rays: general", "gravitational waves", "methods: observational", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1983NIMPR.212..319H", "1989Natur.340..126E", "1992ApJ...395L..83N", "1992Sci...256..325A", "1996ApJ...461..396M", "1998PhRvD..57.3873F", "2002AJ....123.1086H", "2003A&A...411L.291M", "2005ApJ...622..759G", "2005SSRv..120..143B", "2007ApJ...657.1026W", "2007NJPh....9...17L", "2007PhR...442..166N", "2009ApJ...697.1071A", "2009ApJ...702..791M", "2009GCN..9334....1O", "2009MNRAS.398.1601F", "2009PhRvD..80j2001A", "2010ApJ...709L.146D", "2010ApJ...719..900K", "2011APh....35..230V", "2011arXiv1103.2987R", "2012APh....35..641A", "2013ApJ...763...71A", "2013ApJ...771...57A", "2013ApJ...779L...1K", "2013GCN.15464....1R", "2013PASP..125..306F", "2013arXiv1303.4054C", "2014ChPhC..38i0001O", "2015ApJ...806...52S", "2015ApJS..218...23A", "2015NIMPA.774..103B", "2015PhRvD..91f2008A", "2015arXiv150203122V", "2016ApJ...819L..21L", "2016ApJ...820L..36S", "2016ApJ...821L..18P", "2016ApJ...823L...2A", "2016ApJ...826L...6C", "2016ApJ...826L..13A", "2016ApJ...827L..38G", "2016ApJS..223...26A", "2016GCN.19403....1L", "2016GCN.19580....1D", "2016MNRAS.455..712L", "2016PhRvL.116f1102A", "2016PhRvL.116x1103A", "2016arXiv160207352L", "2017ApJ...835...82R", "2017NewA...51....7J", "2018JCAP...04..054F" ]
[ "10.3847/2041-8213/aa7262", "10.48550/arXiv.1607.01793" ]
1607
1607.01793_arXiv.txt
The first direct detection of a Gravitational Wave event (GW) by the recently upgraded \LIGO \citep{1992Sci...256..325A,2009PhRvD..80j2001A, 2016PhRvL.116f1102A} opened a new era in astronomy. The first science run `O1' with the Advanced \LIGO detector started in September 2015, and two high-significance events (GW150914 and GW151226) and one sub-threshold event (LVT151012) were reported \citep{2016PhRvL.116f1102A, 2016PhRvL.116x1103A}. These three events were compatible with the signal expected from the merger of two Black Holes (BH). In future \LIGO science runs, additional sources of GW events might include the mergers of other compact object binary systems: neutron star-black hole (NS-BH) and neutron star-neutron star (NS-NS). The identification and study of electromagnetic counterparts (EM) to GW events is important for several reasons. It resolves degeneracies associated with the inferred binary parameters and allows for a cross-check between the distances measured through the GW signal with the redshifts measured through the EM counterpart, providing an independent constraint on cosmological models. The simultaneous detection of a clear EM counterpart can also confirm near-threshold or sub-threshold GW events, effectively increasing the sensitivity of the search and the distance to which GW events can be detected by \LIGO/\VIRGO. The potential wealth of complementary information encoded in the EM signal is likewise essential to fully unravel the astrophysical context of the event. However, discovering an EM counterpart is challenging because localization regions of GW event provided by \LIGO/\VIRGO are currently as large as several hundred square degrees, much larger than the fields of view (FoVs) of typical X-ray, optical or radio telescopes \citep{2016arXiv160208492A}. The luminosity of the expected EM counterpart is also expected to decay rapidly, leaving a short time window to complete the coverage of the localization region. On the other hand, hard X-ray telescopes such as \Swift-BAT \citep{2005SSRv..120..143B} and \INTEGRAL-ISGRI \citep{2003A&A...411L.291M}, as well as $\gamma$-ray detectors such as the \Fermi Gamma-Ray Burst Monitor \citep[GBM,][]{2009ApJ...702..791M} \Fermi Large Area Telescope \citep[LAT,][]{2009ApJ...697.1071A} and HAWC \citep{2012APh....35..641A}, have much larger FoVs and can cover the localization region much more quickly. They are therefore expected to play a major role in the discovery of the first EM counterpart to a GW event. Short gamma-ray bursts (sGRBs) are a class of GRB with durations $\lesssim 2$ s, and they are thought to be associated with the mergers of BH-NS or NS-NS binaries (\citealt{Eichler1989}, \citealt{Narayan1992}, \citealt{LeeRamirezRuiz2007}, and \citealt{Nakar2007}). They are therefore the expected EM counterparts for GW events involving at least one NS. \Fermi-GBM is the most prolific detector of sGRBs ($\sim 40$ per year), and it is likely to be the first instrument to firmly detect an EM counterpart to a GW event \citep{GW150914_GBM}. However, it localizes sGRBs with uncertainties of the order of a few degrees, making the follow-up by instruments at other wavelengths challenging. Coded-mask telescopes such as BAT and ISGRI can localize events with arcmin precision, but they have smaller FoVs and indeed detect $\sim 8$ sGRBs per year. HAWC has a very large FoV, but as yet has not detected its first GRB. The \Fermi observatory was launched in June 2008 and orbits the Earth at an altitude of $\sim 560$ km with a period of 96.5 minutes. The \Fermi-LAT is a pair-conversion telescope that detects $\gamma$ rays in the energy range from 20\,MeV to more than 300\,GeV. It has a FoV of $\sim 2.4$ sr and it covers the entire sky every $\sim 3$ hours during normal operations. It detects around 15 GRBs per year, among them 2--3 are of the short-duration class, with localization of the order of $\sim 10$ arcmin \citep{2015arXiv150203122V}. When detected by the LAT at high energy ($>$ 100 MeV), sGRBs have a substantially longer duration with respect to their keV--MeV emission. This long-lasting emission is thought to be related to the so-called afterglow emission, also observed at other wavelengths \citep{grb090510,grb110731,kouveliotou13}. \Fermi-LAT is the only wide-field instrument that has detected and localized an sGRB \textit{during its afterglow phase} starting from the GBM localization. MASTER and iPTF have been able to do the same, but only for \textit{long} GRBs so far \citet{2016MNRAS.455..712L, 2015ApJ...806...52S}. Should the detection of an EM counterpart be made by the GBM, \Fermi-LAT could substantially reduce the localization uncertainty, facilitating follow-up at other wavelengths. Should the counterpart be occulted by the Earth for the GBM, and outside the FoV of coded mask instruments, then the LAT would be the only instrument that could still detect the GRB in the $1$--$2$ hours after the burst. Therefore, \Fermi-LAT plays a unique role in facilitating the multi-wavelength follow-up of GW events. This happens routinely already for GRBs, as the vast majority of GRBs detected by the GBM that were also localized by the LAT, were then successfully followed up by other instruments. Among other results, this led to the spectroscopic measurements of more than 20 redshifts \citep{2015arXiv150203122V} for GRBs. Very bright GRBs can also be localized by the LAT on-board with an accuracy between $0.1$ and $\sim 0.5$ deg. This localization is then distributed within seconds to observatories on the ground, allowing for quick follow-up. This happened four times during the first 8 years of the \Fermi mission \citep{2009GCN..9334....1O,2013GCN..15464...1R,2016GCN..19403....1O,2016GCN..19580....1O}. \blue{BH-BH mergers are sources of GW as well, but are not expected to produce EM signals. Nonetheless the \Fermi-GBM observation of a low-significance candidate counterpart 0.4 s after GW 150914 warrants searches for counterparts for GW produced by this class of progenitors.}\footnote{We note that \citet{Lyutikov2016} contests this association on the grounds of the constraints it imposes on the circum-merger environment, while other authors argue against it mainly due to the non-detection by the \INTEGRAL-ACS instrument \citep{greiner16,GW150914_integral}.} The possible association between BH-BH mergers and $\gamma$-ray transients will be addressed by future GW events. If confirmed, it would constitute a surprise that would require new models, such as those published following the report of GW150914-GBM \citep[e.g.][]{Loeb2016, Fraschetti2016, Janiuk2016, 2016ApJ...821L..18P}. Some of these new models foresee a counterpart similar to a standard sGRB, that would imply a possible afterglow signal in the LAT. The standard LAT analysis assumes the source location to be known with some accuracy. Given the size of the localization region of a GW event, the search for a transient counterpart in LAT data is challenging and requires \emph{ad-hoc} methods. In the case of a non-detection, constraining the flux of the EM counterpart requires accounting for the uncertainty on the position of the source, which requires a careful statistical treatment. In this paper, we detail two new methods to search for EM counterparts to GW events in \Fermi-LAT data and to constrain their fluxes. A comprehensive presentation of the results of the \Fermi-LAT follow up for three GW events using the methods presented here is provided in \citet{GW150914_LAT} and \citet{Racusin16}.
\label{sec:disc} With its wide FoV and its survey capabilities, \Fermi-LAT is suitable for looking for EM counterparts to GW events above 100\,MeV and for constraining their fluxes. In particular, in cases of NS-NS or NS-BH mergers, the expected EM counterpart is an sGRB. The LAT is a wide-field instrument that routinely detects and localizes GRBs during their afterglow phase. If the position measured by other instruments such as \Fermi-GBM during the prompt phase has an uncertainty too large for follow up by X-ray, optical and radio telescopes, a LAT detection and localization can vastly improve the chances for a successful follow up. Moreover, the LAT is one of few instruments that can constrain the flux from the EM counterpart despite a large uncertainty on its position. We have presented two novel techniques to perform the search for an EM counterpart to a GW event in \Fermi-LAT data and to provide constraints on its flux. They fully exploit at the same time the capabilities of the instrument as well as the prior information available from the \LIGO/\VIRGO observatories. These methods, developed during the first \LIGO science run `O1', will be systematically used to search for EM counterparts to future GW events. In case of a detection, the methods presented here will return a localization, a flux estimation and a significance of the EM counterpart. If no EM counterpart is detected, a meaningful set of constraints on the flux of the source can be measured. \blue{The methods are now implemented in an automatic analysis pipeline triggered by \LIGO/\VIRGO. This minimizes the turn-around time and increases the chance, in case of a detection, of initiating a prompt follow-up campaign to detect the EM counterpart at other wavelengths.} The authors thank A.~Strong (\blue{Max-Planck institute f\"ur extraterrestrische Physik}), J.~Conrad (Stockholm University) and N.M.~Mazziotta (\blue{INFN Sezione di Bari}) for the helpful discussion on the STAT mailing list of the \Fermi-LAT collaboration. Some of the results in this paper have been derived using the HEALPix \citep{HEALPix} package. The \Fermi-LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat \`a l'Energie Atomique and the Centre National de la Recherche Scientifique / Institut National de Physique Nucl\'eaire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK) and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K.~A.~Wallenberg Foundation, the Swedish Research Council and the Swedish National Space Board in Sweden. Additional support for science analysis during the operations phase is gratefully acknowledged from the Istituto Nazionale di Astrofisica in Italy and the Centre National d'\'Etudes Spatiales in France.
16
7
1607.01793
At least a fraction of gravitational-wave (GW) progenitors are expected to emit an electromagnetic (EM) signal in the form of a short gamma-ray burst (sGRB). Discovering such a transient EM counterpart is challenging because the LIGO/VIRGO localization region is much larger (several hundreds of square degrees) than the field of view of X-ray, optical, and radio telescopes. The Fermi Large Area Telescope (LAT) has a wide field of view (∼2.4 sr) and detects ∼2-3 sGRBs per year above 100 MeV. It can detect them not only during the short prompt phase, but also during their long-lasting high-energy afterglow phase. If other wide-field, high-energy instruments such as Fermi-GBM, Swift-BAT, or INTEGRAL-ISGRI cannot detect or localize with enough precision an EM counterpart during the prompt phase, the LAT can potentially pinpoint it with ≲ 10 arcmin accuracy during the afterglow phase. This routinely happens with gamma-ray bursts. Moreover, the LAT will cover the entire localization region within hours of any triggers during normal operations, allowing the γ-ray flux of any EM counterpart to be measured or constrained. We illustrate two new ad hoc methods to search for EM counterparts with the LAT and their application to the GW candidate LVT151012.
false
[ "EM counterparts", "EM", "radio telescopes", "LAT", "the short prompt phase", "the afterglow phase", "enough precision", "the prompt phase", "square degrees", "≲", "normal operations", "an EM counterpart", "any EM counterpart", "view", "gamma-ray bursts", "LVT151012", "∼2.4 sr", "GW", "≲ 10 arcmin accuracy", "their long-lasting high-energy afterglow phase" ]
15.973012
1.288952
5
2074333
[ "Johnson, Bradley R.", "Flanigan, Daniel", "Abitbol, Maximilian H.", "Ade, Peter A. R.", "Bryan, Sean", "Cho, Hsiao-Mei", "Datta, Rahul", "Day, Peter", "Doyle, Simon", "Irwin, Kent", "Jones, Glenn", "Kernasovskiy, Sarah", "Li, Dale", "Mauskopf, Philip", "McCarrick, Heather", "McMahon, Jeff", "Miller, Amber", "Pisano, Giampaolo", "Song, Yanru", "Surdi, Harshad", "Tucker, Carole" ]
2016SPIE.9914E..0XJ
[ "Polarization sensitive Multi-Chroic MKIDs" ]
9
[ "Columbia Univ. (United States)", "Columbia Univ. (United States)", "Columbia Univ. (United States)", "Cardiff Univ. (United Kingdom)", "Arizona State Univ. (United States)", "SLAC National Accelerator Lab. (United States)", "Univ. of Michigan (United States)", "Jet Propulsion Lab. (United States)", "Cardiff Univ. (United Kingdom)", "Stanford Univ. (United States)", "Columbia Univ. (United States)", "Stanford Univ. (United States)", "SLAC National Accelerator Lab. (United States)", "Arizona State Univ. (United States)", "Columbia Univ. (United States)", "Univ. of Michigan (United States)", "Columbia Univ. (United States)", "Cardiff Univ. (United Kingdom)", "Stanford Univ. (United States)", "Arizona State Univ. (United States)", "Cardiff Univ. (United Kingdom)" ]
[ "2017ApPhL.110v2601J", "2017arXiv170602464A", "2018JLTP..193..103J", "2018JLTP..193..149T", "2018JLTP..193..170T", "2018JLTP..193..633H", "2020JLTP..199..362T", "2020JLTP..200..353K", "2022arXiv220315902B" ]
[ "astronomy", "physics" ]
14
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "2004ApPhL..85.2137D", "2006NIMPA.559..561D", "2008ApPhL..92u2504G", "2008JLTP..151..530D", "2008JLTP..151..684S", "2010ITAP...58.1383Z", "2010SPIE.7741E..1SN", "2010SPIE.7741E..1VD", "2011ApJS..194...24M", "2012JLTP..167..852S", "2012SPIE.8452E..05G", "2012SPIE.8452E..1CK", "2012SPIE.8452E..1DA", "2012SPIE.8452E..2GK", "2013ITAS...2300404P", "2014ApJ...792...62B", "2014JAI.....340001G", "2014JLTP..176..670D", "2014RScI...85l3117M", "2014SPIE.9153E..1PB", "2016PhDT........51S" ]
[ "10.1117/12.2233243", "10.48550/arXiv.1607.03796" ]
1607
1607.03796_arXiv.txt
\label{sec:introduction} Microwave kinetic inductance detectors (MKIDs) are superconducting thin-film, GHz resonators that are designed to also be optimal photon absorbers\cite{zmu}. Absorbed photons with energies greater than the superconducting gap ($\nu > 2 \Delta/h \cong 74~\mbox{GHz} \times (T_c/1~\mbox{K})$) break Cooper pairs, changing the density of quasiparticles in the device. The quasiparticle density affects the kinetic inductance and the dissipation of the superconducting film, so a changing optical signal will cause the resonant frequency and internal quality factor of the resonator to shift. These changes in the properties of the resonator can be detected as changes in the amplitude and phase of a probe tone that drives the resonator at its resonant frequency. This detector technology is particularly well-suited for sub-kelvin, kilo-pixel detector arrays because each detector element can be dimensioned to have a unique resonant frequency, and the probe tones for hundreds to thousands of detectors can be carried into and out of the cryostat on a single pair of coaxial cables. In this paper, we report on the development of modular arrays of horn-coupled, polarization-sensitive MKIDs that are each sensitive to two spectral bands between 125 and 280~GHz. The scalable prototype MKID arrays we are developing are tailored for future multi-kilo-pixel experiments that are designed to simultaneously characterize the polarization properties of both the cosmic microwave background (CMB) and Galactic dust emission. Our device design builds from successful transition edge sensor (TES) bolometer architectures that have been developed by the Truce Collaboration\footnote{http://casa.colorado.edu/$\sim$henninjw/TRUCE/TRUCE.html} and demonstrated to work in receivers on the ACT and SPT telescopes\cite{thornton2016,actpol,austermann2012,sptpol}. Detector modules like these could be a strong candidate for a future CMB satellite mission and/or CMB-S4\cite{bock_2009,cmbs4} because these future multi-kilo-pixel programs will require efficient multiplexing schemes and MKID arrays could out-perform current technologies in this regard (see Figure~\ref{fig:focal_plane_concept}). \begin{figure}[!t] \centering \includegraphics[width=0.95\textwidth]{./focal_plane_development_v3.pdf} \caption{ \textbf{Left:} A scalable 20-element prototype module for dual-polarization LEKIDs. The module we are currently building for our prototype multi-chroic MKID arrays looks similar. \textbf{Right:} Future focal plane concept. By design, the module architecture shown on the left is scalable to one of the seven modules shown on the right. The concept detector array on the right includes 9268 single polarization detectors spread over two spectral bands. } \label{fig:focal_plane_concept} \end{figure} A range of MKID-based instruments have already shown that MKIDs work well at millimeter and sub-millimeter wavelengths. Early MKIDs used antenna coupling\cite{Day2006}, and these antenna-coupled MKIDs were demonstrated at the Caltech Submillimeter Observatory (CSO) in 2007\cite{Schlaerth2008} leading to the development of MUSIC, a multi-chroic antenna-coupled MKID camera\cite{golwala+12}. A simpler device design that uses the inductor in a single-layer $LC$ resonator to directly absorb the millimeter and sub-millimeter-wave radiation was published in 2008\cite{doyle}. This style of MKID, called the lumped-element kinetic inductance detector (LEKID), was first demonstrated in 2011 in the 224-pixel NIKA dual-band millimeter-wave camera on the IRAM 30~m telescope in Spain\cite{monfardini}. Laboratory studies have shown that state-of-the-art MKID and LEKID designs can achieve photon noise limited performance\cite{mauskopf14,McKenney2012}. Photon noise limited horn-coupled LEKIDs sensitive to 1.2~THz were recently demonstrated\cite{Hubmayr2014} and these detectors will be used in BLAST-TNG\cite{galitzki2014,dober2014}. Members of our collaboration conducted studies of horn-coupled LEKIDs to see if LEKIDs would be suitable for CMB polarimetry. These studies revealed that the sensitivity of the LEKID variety of MKID can be compared with that of state-of-the-art TES bolometers\cite{flanigan_2016,mccarrick_2014}. On-chip spectrometers based on MKIDs are currently being developed\cite{superspec2012,microspec2013}. And a large format sub-millimeter wavelength camera, called A-MKID, with more than 10,000 pixels and readout multiplexing factors greater than 1,000 has been built and is currently being commissioned\cite{baryshev+2014}. \begin{figure}[!t] \centering \includegraphics[width=0.9\textwidth]{./project.pdf} \caption{ \textbf{Top~Left:} A cross-sectional view of one focal-plane element. In an effort to minimize two-level system (TLS) noise, the MKID sensing element is deposited directly on the silicon wafer, and it is not covered with silicon nitride. \textbf{Bottom~Left:} A scale drawing of the dual-polarization multi-chroic MKID device we propose to develop. \textbf{Bottom~Right:} A schematic of the co-planar waveguide MKID we are developing. Photons from the sky are brought to the detector on a microstrip from the hybrid tee. These photons then couple to our resonator and are absorbed in the aluminum section of the CPW $\lambda/4$ resonator. \textbf{Top~Right:} End-to-end electromagnetic simulations show the expected absorption efficiency is approximately 90\% across the 150~GHz and the 235~GHz spectral bands. } \label{fig:pixel_design} \end{figure}
\label{sec:discussion} One of the primary goals of this project is to bring the functionality of MKIDs for CMB studies in line with state-of-the-art TES bolometers\cite{bicep2_inst_2014,suzuki2012,arnold2012,polarbear12,rostem_2014,benson2014}. To date we have (i) designed the critical broadband microstrip-to-CPW coupler, (ii) modified the existing Advanced ACTPol design for SOI, (iii) fabricated MKID optimization chips, and (iv) developed several of the critical fabrication steps. We are currently dark testing the MKIDs and laying out our final array design. Prototype array fabrication will begin in the summer of 2016. In the future, we plan on making the sensing element in the MKIDs out of aluminum manganese instead of aluminum. By adding manganese to the aluminum, the $T_c$ of the sensor decreases in a controllable way\cite{deiker_2004}, which does two critical things. First and foremost, in our current 150 and 235~GHz spectral bands, the photons are energetic enough to break multiple Cooper pairs in the sensing element, so the detector noise will be suppressed below the photon noise -- even for the low optical loads that are expected in a space-like environment. Second, a lower $T_c$ makes the detector technology sensitive to lower frequencies, so this technology will open the door to low-frequency ($\sim$30~GHz) MKIDs in the future.
16
7
1607.03796
We report on the development of scalable prototype microwave kinetic inductance detector (MKID) arrays tai- lored for future multi-kilo-pixel experiments that are designed to simultaneously characterize the polarization properties of both the cosmic microwave background (CMB) and Galactic dust emission. These modular arrays are composed of horn-coupled, polarization-sensitive MKIDs, and each pixel has four detectors: two polariza- tions in two spectral bands between 125 and 280 GHz. A horn is used to feed each array element, and a planar orthomode transducer, composed of two waveguide probe pairs, separates the incoming light into two linear po- larizations. Diplexers composed of resonant-stub band-pass filters separate the radiation into 125 to 170 GHz and 190 to 280 GHz pass bands. The millimeter-wave power is ultimately coupled to a hybrid co-planar waveguide microwave kinetic inductance detector using a novel, broadband circuit developed by our collaboration. Elec- tromagnetic simulations show the expected absorption efficiency of the detector is approximately 90%. Array fabrication will begin in the summer of 2016.
false
[ "scalable prototype microwave kinetic inductance detector", "Galactic dust emission", "a hybrid co-planar waveguide microwave kinetic inductance detector", "resonant-stub band-pass filters", "Array fabrication", "future multi-kilo-pixel experiments", "two spectral bands", "190 to 280 GHz pass bands", "170 GHz", "280 GHz", "CMB", "Galactic", "MKID", "two waveguide probe pairs", "the polarization properties", "both the cosmic microwave background", "the incoming light", "four detectors", "the detector", "a planar orthomode transducer" ]
10.337198
3.0633
68
12406030
[ "Juvela, M." ]
2016A&A...593A..58J
[ "Template matching method for the analysis of interstellar cloud structure" ]
10
[ "Department of Physics, PO Box 64, University of Helsinki, 00014, Helsinki, Finland ; Institut UTINAM, UMR 6213, CNRS, Univ. Bourgogne Franche Comte, 41bis avenue de l'Observatoire, 25000, Besançon, France" ]
[ "2018A&A...614A..83J", "2021A&A...649A..89M", "2021A&A...654A..78M", "2022A&A...668A..41C", "2022IEEEA..1074472A", "2022arXiv220500683A", "2023ASPC..534..233P", "2023MNRAS.518.5904Y", "2023MNRAS.524.2994H", "2024RvMPP...8...21Y" ]
[ "astronomy" ]
3
[ "ISM: clouds", "ISM: structure", "infrared: ISM", "submillimeter: ISM", "techniques: image processing", "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
[ "1981Natur.291..561A", "1990ASSL..162..151S", "1990ApJ...356..513S", "1991ApJ...368..432L", "1991ApJ...376..561M", "1992AJ....103.1313G", "1993MNRAS.262..327G", "1996A&A...314..625A", "1997ApJ...486..944E", "1998A&A...336..150M", "1999MNRAS.305..143W", "2002A&A...384..225S", "2002ApJ...580L..57P", "2003A&A...398..785S", "2003ApJ...583..308P", "2009A&A...508L..35K", "2010A&A...518L...1P", "2010A&A...518L.106K", "2011A&A...529L...6A", "2011A&A...530A.133M", "2011MNRAS.414..350S", "2012A&A...541A..12J", "2012A&A...544A.141J", "2013A&A...556A.153H", "2013A&A...560A..63M", "2014ApJ...789...82C", "2014ApJ...791...27S", "2014prpl.conf...27A", "2015A&A...584A..92M", "2015A&A...584A..93J", "2015A&A...584A..94J", "2015MNRAS.449.1782S", "2015MNRAS.453L..41S", "2015NatSR...518267B", "2016A&A...591A..90R", "2016MNRAS.460.1934M" ]
[ "10.1051/0004-6361/201628727", "10.48550/arXiv.1607.01931" ]
1607
1607.01931_arXiv.txt
\label{sect:intro} The characterisation of the structure of interstellar medium (ISM) is difficult because of its complexity. The structures do not generally follow any simple patterns that could be easily parameterised. The overall properties of interstellar medium can be examined using general statistics such as the distribution of column density values ($P(D)$ analysis) \citep[e.g.][]{Kainulainen2009, Schneider2015}, power spectra and structure functions \citep[e.g.][]{Armstrong1981,Green1993,Gautier1992,Padoan2002,Padoan2003}, or fractal analysis \citep[e.g.][]{Scalo1990, Elmegreen1997}. These methods do not directly provide information on the shapes of individual objects, only on the global statistics of a field. Regarding individual structures, the analysis has concentrated on the smallest scales, on individual clumps and pre-stellar and protostellar cores \citep[e.g.][]{Andre1996,Motte1998,WardThompson1999,Schneider2002,Konyves2010}. Because of the approximate balance of gravity and supporting forces, the shapes at the smallest scales may be approximated with one-dimensional density profiles or simple 2D shapes such as 2D Gaussians \citep[e.g.][]{Stutzki1990,Lada1991,Myers1991}, although the apparent simplicity is sometimes related to the finite resolution of observations. Recently the interest in the intermediate, better resolved scales has increased. The interstellar cloud filaments imaged by the {\it Herschel Space Observatory} \citep{Pilbratt2010} are one example of this. In the nearby star-forming regions even the internal structure of filaments can be resolved \citep{Arzoumanian2011, Juvela2012_filsimu} and detection of filamentary structures is possible up to kiloparsec distances. In the context of interstellar clouds, the importance of filaments is partly connected to the role they play in star formation \citep{Andre2014PPVI}. However, filamentary structures can be detected in almost any kind of astronomical sources, from the Sun to the distant universe, in both continuum and line data. Several methods have been used to identify and to characterise elongated structures from 3D data (usually simulations), line observations (position-position-velocity data), and two-dimensional images. An incomplete list of methods includes DisPerSe \citep{Sousbie2011, Arzoumanian2011} and getfilaments \citep{Menshchikov2013, Rivera2016a} routines, derivative-based (Hessian) methods \citep{Molinari2011, Schisano2014, Salji2015}, curvelets and ridgelets \citep[e.g.][]{Starck2003}, and the inertia matrix method \citep{Hennebelle2013}. The tools vary regarding their complexity, computational cost, and the nature and detail of the information they return on the identified structures. In this paper, we examine the performance of a simple template matching (TM) technique in the structural analysis, for detecting for example the presence of elongated (filamentary) features, in two-dimensional image data. We compare the method to rolling Hough transform (RHT), which has been used in many image analysis applications \citep{Illingworth1988} and has been applied to the filament detection from H{\rm I} observations \citep{Clark2014} and from $Herschel$ continuum data \citep{Malinen2016}. RHT was selected because its basic principles are very similar to the proposed TM implementation. Both methods are used to enhance certain structural features of the input images and thus to help characterise the structures, for example, regarding the degree of angular anisotropy. Neither method is geared towards the extraction of individual sources. Our aim is to show that TM provides a fast and yet relatively robust way to identify local anisotropy and that, with minor changes, can also be used to highlight other types of structures. The content of this paper is as follows. Section~\ref{sect:methods} gives an overview of the TM and RHT methods. Further details on our TM implementation are given in Appendices~\ref{appendix:implementation} and \ref{appendix:alternative}. In Sect.~\ref{sect:results} we present a series of tests where the results of RHT and our TM implementation are compared for synthetic images. We discuss the results and possible applications of the TM method in Sect.~\ref{sect:discussion} before summarising our conclusions in Sect.~\ref{sect:conclusions}. \footnote{Our python implementation of the TM routine can be obtained from {\tt www.interstellarmedium.org/PatternMatching}.}
\begin{itemize} \item Template matching techniques provide a simple and yet an efficient method of image analysis in terms of both computational cost and the quality of results. \item In detection of elongated structures, TM provides results comparable to those of RHT. However, TM was found to perform better when maps contain significant noise and background fluctuations. \item Small changes of the basic algorithm (e.g. regarding data normalisation and spatial filtering) can be used to emphasise the brightest structures or to highlight structures at very faint levels. \item In the case of $Herschel$, its data quality allows the identification of filamentary structures at very low surface brightness levels, down to column densities $N({\rm H}_2) \sim 10^{20}$\,cm$^{-2}$. \end{itemize}
16
7
1607.01931
Context. The structure of interstellar medium can be characterised at large scales in terms of its global statistics (e.g. power spectra) and at small scales by the properties of individual cores. Interest has been increasing in structures at intermediate scales, resulting in a number of methods being developed for the analysis of filamentary structures. <BR /> Aims: We describe the application of the generic template-matching (TM) method to the analysis of maps. Our aim is to show that it provides a fast and still relatively robust way to identify elongated structures or other image features. <BR /> Methods: We present the implementation of a TM algorithm for map analysis. The results are compared against rolling Hough transform (RHT), one of the methods previously used to identify filamentary structures. We illustrate the method by applying it to Herschel surface brightness data. <BR /> Results: The performance of the TM method is found to be comparable to that of RHT but TM appears to be more robust regarding the input parameters, for example, those related to the selected spatial scales. Small modifications of TM enable one to target structures at different size and intensity levels. In addition to elongated features, we demonstrate the possibility of using TM to also identify other types of structures. <BR /> Conclusions: The TM method is a viable tool for data quality control, exploratory data analysis, and even quantitative analysis of structures in image data.
false
[ "exploratory data analysis", "filamentary structures", "structures", "elongated structures", "map analysis", "image data", "data quality control", "other image features", "small scales", "large scales", "intermediate scales", "Herschel surface brightness data", "individual cores", "TM", "elongated features", "methods", "other types", "even quantitative analysis", "the selected spatial scales", "The structure" ]
11.571614
10.080662
183
12406155
[ "Klochkov, D.", "Suleimanov, V.", "Sasaki, M.", "Santangelo, A." ]
2016A&A...592L..12K
[ "Study of a new central compact object: The neutron star in the supernova remnant G15.9+0.2" ]
23
[ "Institut für Astronomie und Astrophysik (IAAT), Universität Tübingen, Sand 1, 72076, Tübingen, Germany", "Institut für Astronomie und Astrophysik (IAAT), Universität Tübingen, Sand 1, 72076, Tübingen, Germany", "Institut für Astronomie und Astrophysik (IAAT), Universität Tübingen, Sand 1, 72076, Tübingen, Germany", "Institut für Astronomie und Astrophysik (IAAT), Universität Tübingen, Sand 1, 72076, Tübingen, Germany" ]
[ "2017A&A...597A..65M", "2017A&A...600A..43S", "2017ApJ...844...84H", "2017JPhCS.932a2006D", "2018A&A...609A..74P", "2018A&A...618A..76D", "2018MNRAS.479.3033S", "2019ApJ...875..122A", "2019MNRAS.483...70C", "2019MNRAS.484..974W", "2019PASP..131k4301T", "2019PhRvC.100e5801B", "2020ApJ...904...70L", "2020MNRAS.491.1585H", "2020MNRAS.493.2706C", "2020MNRAS.496.5052P", "2021ARep...65..615K", "2021MNRAS.506..709S", "2021MNRAS.506.5015H", "2021PhR...919....1B", "2021RAA....21..294W", "2023A&A...673A..15S", "2023ApJ...944...36A" ]
[ "astronomy" ]
5
[ "stars: neutron", "ISM: supernova remnants", "stars: atmospheres", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1979ApJ...228..939C", "1982GeCoA..46.2363A", "1983ApJ...270..119M", "1987PASA....7..167P", "2002ASPC..271..247P", "2004ApJS..155..623P", "2004IAUS..218..239P", "2006ApJ...652L..45R", "2006NuPhA.777..497P", "2006SPIE.6270E..1VF", "2007ASSL..326.....H", "2009ApJ...703..910P", "2009Natur.462...71H", "2010ApJ...709..436H", "2013A&A...556A..41K", "2013ApJ...765...58G", "2014ApJ...790...94B", "2014ApJS..210...13S", "2015A&A...573A..53K", "2015MNRAS.452..540B", "2015MNRAS.454.2668O", "2016PhR...621..127L", "2016RvMP...88b1001W" ]
[ "10.1051/0004-6361/201629208", "10.48550/arXiv.1607.08021" ]
1607
1607.08021_arXiv.txt
Neutron stars are fundamental objects of modern astrophysics because (i) supranuclear densities in their interiors probably yield unexplored forms of matter, (ii) they are the outcome of the explosive collapse of massive stars, and (iii) their extreme gravitational and magnetic fields give rise to reach phenomenology of the associated emission and accretion. Only a few dozen neutron stars (NSs) are associated with their parent supernova remnants (SNRs). A dozen of these stars show pure thermal X-ray emission and lack any magnetic activity such as flaring, nonthermal magnetospheric emission or pulsar wind nebulae. This subgroup are usually referred to as central compact objects (CCOs) and were first considered as a separate class of isolated NSs by \citet{Pavlov:etal:02,Pavlov:etal:04} based on \emph{Chandra} observations. For three CCOs, pulsations with periods in the range of $\sim$0.1--0.4\,s have been detected. The measured spin-down rates confirm relatively low magnetic fields of these objects, namely, $B\sim 3\times 10^{10}- 10^{11}$\,G \citep{Halpern:Gotthelf:10,Gotthelf:etal:13}, which is substantially weaker than in radio and accreting pulsars. The class of CCOs is thus believed to be composed by young (a few 10$^3-10^4$\,yr), weakly magnetized, thermally emitting, cooling NSs. One of the major efforts in NS studies is the modeling of the thermal emission from the stellar surface. Under certain assumptions, such a modeling permits constraints on the geometrical and physical properties of the stars and, most importantly, on their mass, radius, and effective temperature. Constraints on the mass and radius allow one to probe the equation of state of the superdense matter in the stellar interior \citep[e.g.,][]{Haensel:etal:07,Lattimer:Prakash:15,Watts:etal:16}. The effective temperature of the stellar surface combined with the estimated age of the object allows one to probe the NS cooling rate \citep[e.g.,][]{Page:etal:04,Page:etal:06}. The cooling mechanisms also strongly depend on the microphysics inside the star \citep[e.g.,][]{Ofengeim:etal:15,Beznogov:Yakovlev:15}. The relatively weak magnetic fields of CCOs permit modeling of their thermal emission under simplified assumptions, namely neglecting the effects of strong magnetic fields in the stellar atmosphere \citep[e.g.,][and references therein]{Suleimanov:etal:14}; this makes these objects perfect laboratories for the study of neutron star physics. The point source \object{CXOU\,J181852.0$-$150213} (hereafter, CXOU\,J1818) was serendipitously discovered in the \emph{Chandra} X-ray observations of the radio-bright, few thousand year old SNR \object{G15.9+0.2} \citep{Reynolds:etal:06}. The relatively sparse data did not permit a detailed spectral analysis of the source. \citet{Reynolds:etal:06} argued, however, that CXOU\,J1818 is probably a NS associated with the remnant. A rotation-powered pulsar or a CCO were considered by the authors to be viable possibilities. The object was subsequently listed among ``CCO candidates'' in \citet{Gotthelf:etal:13}. In this Letter, we present our analysis of the new \emph{Chandra} observation of G15.9+0.2 and CXOU\,J1818, which has increased the total exposure time for these sources by a factor of four. Our study is focused on the compact source and is fully consistent with the CCO hypothesis. The detailed analysis of the diffuse emission of the SNR is not part of this work and will be presented elsewhere.
Our spectral analysis of the central source in the SNR G15.9+0.2 clearly supports the hypothesis of a thermally emitting NS. The object thus can be considered to be a new member of the CCO class. A power-law spectral model yields an absorption column density that is incompatible with that of the remnant and a photon index that is much softer than expected in case of an AGN, a young radio pulsar, or a pulsar wind nebula. We thus consider these alternatives as rather improbable. We also confirmed the absence of an optical or infrared counterpart with the improved sky coordinates of the source. The stellar coronal origin of the object is, therefore, also unlikely. Similar to other CCOs, the distance to CXOU\,J1818 derived from the blackbody normalization is unrealistically high, $\sim$200\,kpc, assuming that the radiation is emitted by the entire stellar surface. A fit with the hydrogen atmosphere models reduces the distance to $\sim$40--70\,kpc, which is still too high for a Galactic source. Only the carbon atmosphere model yields a meaningful distance of $\sim$10--20\,kpc. Very similar results have been obtained for two other CCOs for which much better observational data are available (higher flux, longer and multiple observations), specifically, the central NSs in SNRs Cas~A \citep{Ho:Heinke:09} and HESS\,J1731$-$347 \citep{Klochkov:etal:13,Klochkov:etal:15}. In these two objects, the hypothesis that the radiation is uniformly emitted by the entire stellar surface (i.e., that the emitting radius is equal to the stellar radius) is supported by the relatively stringent limits on the pulsed fraction of $\lesssim$10\%. Our data do not permit any useful constraints on pulsations of CXOU\,J1818. The pulse periods of the known pulsations CCOs, RX\,J0822.0$-$4300, 1E\,1207.4$-$5209, and CXOU\,J185238.6$+$004020, are all below $\sim$0.4\,s, which is in the range inaccessible with the \emph{Chandra} observations due to insufficient timing resolution. Nevertheless, the similarities in the spectral properties of CXOU\,J1818 with the well-studied CCOs in Cas~A and in HESS\,J1731$-$347 suggest a similar geometry and physics of emission in the three NSs and further supports the association of CXOU\,J1818 with the class of CCOs.
16
7
1607.08021
We present our study of the central point source CXOU J181852.0-150213 in the young Galactic supernova remnant (SNR) G15.9+0.2 based on the recent ~90 ks Chandra observations. The point source was discovered in 2005 in shorter Chandra observations and was hypothesized to be a neutron star associated with the SNR. Our X-ray spectral analysis strongly supports the hypothesis of a thermally emitting neutron star associated with G15.9+0.2. We conclude that the object belongs to the class of young cooling low-magnetized neutron stars referred to as central compact objects (CCOs). We modeled the spectrum of the neutron star with a blackbody spectral function and with our hydrogen and carbon neutron star atmosphere models, assuming that the radiation is uniformly emitted by the entire stellar surface. Under this assumption, only the carbon atmosphere models yield a distance that is compatible with a source located in the Galaxy. In this respect, CXOU J181852.0-150213 is similar to two other well-studied CCOs, the neutron stars in Cas A and in HESS J1731-347, for which carbon atmosphere models were used to reconcile their emission with the known or estimated distances.
false
[ "Chandra observations", "shorter Chandra observations", "central compact objects", "Chandra", "SNR", "Galaxy", "HESS J1731", "Cas A", "a neutron star", "the neutron star", "our hydrogen and carbon neutron star atmosphere models", "which carbon atmosphere models", "G15.9", "CCOs", "young cooling low-magnetized neutron stars", "a thermally emitting neutron star", "the entire stellar surface", "the central point source", "the young Galactic supernova remnant", "only the carbon atmosphere models" ]
4.56506
4.688946
-1
1360289
[ "Abbott, B. P.", "Abbott, R.", "Abbott, T. D.", "Abernathy, M. R.", "Acernese, F.", "Ackley, K.", "Adams, C.", "Adams, T.", "Addesso, P.", "Adhikari, R. X.", "Adya, V. B.", "Affeldt, C.", "Agathos, M.", "Agatsuma, K.", "Aggarwal, N.", "Aguiar, O. D.", "Aiello, L.", "Ain, A.", "Ajith, P.", "Allen, B.", "Allocca, A.", "Altin, P. A.", "Anderson, S. B.", "Anderson, W. G.", "Arai, K.", "Araya, M. C.", "Arceneaux, C. C.", "Areeda, J. S.", "Arnaud, N.", "Arun, K. G.", "Ascenzi, S.", "Ashton, G.", "Ast, M.", "Aston, S. M.", "Astone, P.", "Aufmuth, P.", "Aulbert, C.", "Babak, S.", "Bacon, P.", "Bader, M. K. M.", "Baker, P. T.", "Baldaccini, F.", "Ballardin, G.", "Ballmer, S. W.", "Barayoga, J. C.", "Barclay, S. E.", "Barish, B. C.", "Barker, D.", "Barone, F.", "Barr, B.", "Barsotti, L.", "Barsuglia, M.", "Barta, D.", "Bartlett, J.", "Bartos, I.", "Bassiri, R.", "Basti, A.", "Batch, J. C.", "Baune, C.", "Bavigadda, V.", "Bazzan, M.", "Bejger, M.", "Bell, A. S.", "Berger, B. K.", "Bergmann, G.", "Berry, C. P. L.", "Bersanetti, D.", "Bertolini, A.", "Betzwieser, J.", "Bhagwat, S.", "Bhandare, R.", "Bilenko, I. A.", "Billingsley, G.", "Birch, J.", "Birney, R.", "Biscans, S.", "Bisht, A.", "Bitossi, M.", "Biwer, C.", "Bizouard, M. A.", "Blackburn, J. K.", "Blair, C. D.", "Blair, D. G.", "Blair, R. M.", "Bloemen, S.", "Bock, O.", "Boer, M.", "Bogaert, G.", "Bogan, C.", "Bohe, A.", "Bond, C.", "Bondu, F.", "Bonnand, R.", "Boom, B. A.", "Bork, R.", "Boschi, V.", "Bose, S.", "Bouffanais, Y.", "Bozzi, A.", "Bradaschia, C.", "Brady, P. R.", "Braginsky, V. B.", "Branchesi, M.", "Brau, J. E.", "Briant, T.", "Brillet, A.", "Brinkmann, M.", "Brisson, V.", "Brockill, P.", "Broida, J. E.", "Brooks, A. F.", "Brown, D. A.", "Brown, D. D.", "Brown, N. M.", "Brunett, S.", "Buchanan, C. C.", "Buikema, A.", "Bulik, T.", "Bulten, H. J.", "Buonanno, A.", "Buskulic, D.", "Buy, C.", "Byer, R. L.", "Cabero, M.", "Cadonati, L.", "Cagnoli, G.", "Cahillane, C.", "Calderón Bustillo, J.", "Callister, T.", "Calloni, E.", "Camp, J. B.", "Cannon, K. C.", "Cao, J.", "Capano, C. D.", "Capocasa, E.", "Carbognani, F.", "Caride, S.", "Casanueva Diaz, J.", "Casentini, C.", "Caudill, S.", "Cavaglià, M.", "Cavalier, F.", "Cavalieri, R.", "Cella, G.", "Cepeda, C. B.", "Cerboni Baiardi, L.", "Cerretani, G.", "Cesarini, E.", "Chamberlin, S. J.", "Chan, M.", "Chao, S.", "Charlton, P.", "Chassande-Mottin, E.", "Cheeseboro, B. D.", "Chen, H. Y.", "Chen, Y.", "Cheng, C.", "Chincarini, A.", "Chiummo, A.", "Cho, H. S.", "Cho, M.", "Chow, J. H.", "Christensen, N.", "Chu, Q.", "Chua, S.", "Chung, S.", "Ciani, G.", "Clara, F.", "Clark, J. A.", "Cleva, F.", "Coccia, E.", "Cohadon, P. -F.", "Colla, A.", "Collette, C. G.", "Cominsky, L.", "Constancio., M., Jr.", "Conte, A.", "Conti, L.", "Cook, D.", "Corbitt, T. R.", "Cornish, N.", "Corsi, A.", "Cortese, S.", "Costa, C. A.", "Coughlin, M. 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C.", "Ducrot, M.", "Dwyer, S. E.", "Edo, T. B.", "Edwards, M. C.", "Effler, A.", "Eggenstein, H. -B.", "Ehrens, P.", "Eichholz, J.", "Eikenberry, S. S.", "Engels, W.", "Essick, R. C.", "Etzel, T.", "Evans, M.", "Evans, T. M.", "Everett, R.", "Factourovich, M.", "Fafone, V.", "Fair, H.", "Fairhurst, S.", "Fan, X.", "Fang, Q.", "Farinon, S.", "Farr, B.", "Farr, W. M.", "Favata, M.", "Fays, M.", "Fehrmann, H.", "Fejer, M. M.", "Fenyvesi, E.", "Ferrante, I.", "Ferreira, E. C.", "Ferrini, F.", "Fidecaro, F.", "Fiori, I.", "Fiorucci, D.", "Fisher, R. P.", "Flaminio, R.", "Fletcher, M.", "Fournier, J. -D.", "Frasca, S.", "Frasconi, F.", "Frei, Z.", "Freise, A.", "Frey, R.", "Frey, V.", "Fritschel, P.", "Frolov, V. V.", "Fulda, P.", "Fyffe, M.", "Gabbard, H. A. G.", "Gair, J. R.", "Gammaitoni, L.", "Gaonkar, S. G.", "Garufi, F.", "Gaur, G.", "Gehrels, N.", "Gemme, G.", "Geng, P.", "Genin, E.", "Gennai, A.", "George, J.", "Gergely, L.", "Germain, V.", "Ghosh, Abhirup", "Ghosh, Archisman", "Ghosh, S.", "Giaime, J. A.", "Giardina, K. D.", "Giazotto, A.", "Gill, K.", "Glaefke, A.", "Goetz, E.", "Goetz, R.", "Gondan, L.", "González, G.", "Gonzalez Castro, J. M.", "Gopakumar, A.", "Gordon, N. A.", "Gorodetsky, M. L.", "Gossan, S. E.", "Gosselin, M.", "Gouaty, R.", "Grado, A.", "Graef, C.", "Graff, P. B.", "Granata, M.", "Grant, A.", "Gras, S.", "Gray, C.", "Greco, G.", "Green, A. C.", "Groot, P.", "Grote, H.", "Grunewald, S.", "Guidi, G. M.", "Guo, X.", "Gupta, A.", "Gupta, M. K.", "Gushwa, K. E.", "Gustafson, E. K.", "Gustafson, R.", "Hacker, J. J.", "Hall, B. R.", "Hall, E. D.", "Hammond, G.", "Haney, M.", "Hanke, M. M.", "Hanks, J.", "Hanna, C.", "Hannam, M. D.", "Hanson, J.", "Hardwick, T.", "Harms, J.", "Harry, G. M.", "Harry, I. W.", "Hart, M. J.", "Hartman, M. T.", "Haster, C. -J.", "Haughian, K.", "Heidmann, A.", "Heintze, M. C.", "Heitmann, H.", "Hello, P.", "Hemming, G.", "Hendry, M.", "Heng, I. S.", "Hennig, J.", "Henry, J.", "Heptonstall, A. W.", "Heurs, M.", "Hild, S.", "Hoak, D.", "Hofman, D.", "Holt, K.", "Holz, D. E.", "Hopkins, P.", "Hough, J.", "Houston, E. A.", "Howell, E. J.", "Hu, Y. M.", "Huang, S.", "Huerta, E. A.", "Huet, D.", "Hughey, B.", "Husa, S.", "Huttner, S. H.", "Huynh-Dinh, T.", "Indik, N.", "Ingram, D. R.", "Inta, R.", "Isa, H. N.", "Isac, J. -M.", "Isi, M.", "Isogai, T.", "Iyer, B. R.", "Izumi, K.", "Jacqmin, T.", "Jang, H.", "Jani, K.", "Jaranowski, P.", "Jawahar, S.", "Jian, L.", "Jiménez-Forteza, F.", "Johnson, W. W.", "Jones, D. I.", "Jones, R.", "Jonker, R. J. G.", "Ju, L.", "K, Haris", "Kalaghatgi, C. V.", "Kalogera, V.", "Kandhasamy, S.", "Kang, G.", "Kanner, J. B.", "Kapadia, S. J.", "Karki, S.", "Karvinen, K. S.", "Kasprzack, M.", "Katsavounidis, E.", "Katzman, W.", "Kaufer, S.", "Kaur, T.", "Kawabe, K.", "Kéfélian, F.", "Kehl, M. S.", "Keitel, D.", "Kelley, D. B.", "Kells, W.", "Kennedy, R.", "Key, J. S.", "Khalili, F. Y.", "Khan, I.", "Khan, S.", "Khan, Z.", "Khazanov, E. A.", "Kijbunchoo, N.", "Kim, Chi-Woong", "Kim, Chunglee", "Kim, J.", "Kim, K.", "Kim, N.", "Kim, W.", "Kim, Y. -M.", "Kimbrell, S. J.", "King, E. J.", "King, P. J.", "Kissel, J. S.", "Klein, B.", "Kleybolte, L.", "Klimenko, S.", "Koehlenbeck, S. M.", "Koley, S.", "Kondrashov, V.", "Kontos, A.", "Korobko, M.", "Korth, W. Z.", "Kowalska, I.", "Kozak, D. B.", "Kringel, V.", "Krishnan, B.", "Królak, A.", "Krueger, C.", "Kuehn, G.", "Kumar, P.", "Kumar, R.", "Kuo, L.", "Kutynia, A.", "Lackey, B. D.", "Landry, M.", "Lange, J.", "Lantz, B.", "Lasky, P. D.", "Laxen, M.", "Lazzarini, A.", "Lazzaro, C.", "Leaci, P.", "Leavey, S.", "Lebigot, E. O.", "Lee, C. H.", "Lee, H. K.", "Lee, H. M.", "Lee, K.", "Lenon, A.", "Leonardi, M.", "Leong, J. R.", "Leroy, N.", "Letendre, N.", "Levin, Y.", "Lewis, J. B.", "Li, T. G. F.", "Libson, A.", "Littenberg, T. B.", "Lockerbie, N. A.", "Lombardi, A. L.", "London, L. T.", "Lord, J. E.", "Lorenzini, M.", "Loriette, V.", "Lormand, M.", "Losurdo, G.", "Lough, J. D.", "Lück, H.", "Lundgren, A. P.", "Lynch, R.", "Ma, Y.", "Machenschalk, B.", "MacInnis, M.", "Macleod, D. M.", "Magaña-Sandoval, F.", "Magaña Zertuche, L.", "Magee, R. M.", "Majorana, E.", "Maksimovic, I.", "Malvezzi, V.", "Man, N.", "Mandic, V.", "Mangano, V.", "Mansell, G. L.", "Manske, M.", "Mantovani, M.", "Marchesoni, F.", "Marion, F.", "Márka, S.", "Márka, Z.", "Markosyan, A. S.", "Maros, E.", "Martelli, F.", "Martellini, L.", "Martin, I. W.", "Martynov, D. V.", "Marx, J. N.", "Mason, K.", "Masserot, A.", "Massinger, T. J.", "Masso-Reid, M.", "Mastrogiovanni, S.", "Matichard, F.", "Matone, L.", "Mavalvala, N.", "Mazumder, N.", "McCarthy, R.", "McClelland, D. E.", "McCormick, S.", "McGuire, S. C.", "McIntyre, G.", "McIver, J.", "McManus, D. J.", "McRae, T.", "McWilliams, S. T.", "Meacher, D.", "Meadors, G. D.", "Meidam, J.", "Melatos, A.", "Mendell, G.", "Mercer, R. A.", "Merilh, E. L.", "Merzougui, M.", "Meshkov, S.", "Messenger, C.", "Messick, C.", "Metzdorff, R.", "Meyers, P. M.", "Mezzani, F.", "Miao, H.", "Michel, C.", "Middleton, H.", "Mikhailov, E. E.", "Milano, L.", "Miller, A. L.", "Miller, A.", "Miller, B. B.", "Miller, J.", "Millhouse, M.", "Minenkov, Y.", "Ming, J.", "Mirshekari, S.", "Mishra, C.", "Mitra, S.", "Mitrofanov, V. P.", "Mitselmakher, G.", "Mittleman, R.", "Moggi, A.", "Mohan, M.", "Mohapatra, S. R. P.", "Montani, M.", "Moore, B. C.", "Moore, C. J.", "Moraru, D.", "Moreno, G.", "Morriss, S. R.", "Mossavi, K.", "Mours, B.", "Mow-Lowry, C. M.", "Mueller, G.", "Muir, A. W.", "Mukherjee, Arunava", "Mukherjee, D.", "Mukherjee, S.", "Mukund, N.", "Mullavey, A.", "Munch, J.", "Murphy, D. J.", "Murray, P. G.", "Mytidis, A.", "Nardecchia, I.", "Naticchioni, L.", "Nayak, R. K.", "Nedkova, K.", "Nelemans, G.", "Nelson, T. J. N.", "Neri, M.", "Neunzert, A.", "Newton, G.", "Nguyen, T. T.", "Nielsen, A. B.", "Nissanke, S.", "Nitz, A.", "Nocera, F.", "Nolting, D.", "Normandin, M. E. N.", "Nuttall, L. K.", "Oberling, J.", "Ochsner, E.", "O'Dell, J.", "Oelker, E.", "Ogin, G. H.", "Oh, J. J.", "Oh, S. H.", "Ohme, F.", "Oliver, M.", "Oppermann, P.", "Oram, Richard J.", "O'Reilly, B.", "O'Shaughnessy, R.", "Ottaway, D. J.", "Overmier, H.", "Owen, B. J.", "Pai, A.", "Pai, S. A.", "Palamos, J. R.", "Palashov, O.", "Palomba, C.", "Pal-Singh, A.", "Pan, H.", "Pankow, C.", "Pannarale, F.", "Pant, B. C.", "Paoletti, F.", "Paoli, A.", "Papa, M. A.", "Paris, H. R.", "Parker, W.", "Pascucci, D.", "Pasqualetti, A.", "Passaquieti, R.", "Passuello, D.", "Patricelli, B.", "Patrick, Z.", "Pearlstone, B. L.", "Pedraza, M.", "Pedurand, R.", "Pekowsky, L.", "Pele, A.", "Penn, S.", "Perreca, A.", "Perri, L. M.", "Phelps, M.", "Piccinni, O. J.", "Pichot, M.", "Piergiovanni, F.", "Pierro, V.", "Pillant, G.", "Pinard, L.", "Pinto, I. M.", "Pitkin, M.", "Poe, M.", "Poggiani, R.", "Popolizio, P.", "Post, A.", "Powell, J.", "Prasad, J.", "Predoi, V.", "Prestegard, T.", "Price, L. R.", "Prijatelj, M.", "Principe, M.", "Privitera, S.", "Prix, R.", "Prodi, G. A.", "Prokhorov, L.", "Puncken, O.", "Punturo, M.", "Puppo, P.", "Pürrer, M.", "Qi, H.", "Qin, J.", "Qiu, S.", "Quetschke, V.", "Quintero, E. A.", "Quitzow-James, R.", "Raab, F. J.", "Rabeling, D. S.", "Radkins, H.", "Raffai, P.", "Raja, S.", "Rajan, C.", "Rakhmanov, M.", "Rapagnani, P.", "Raymond, V.", "Razzano, M.", "Re, V.", "Read, J.", "Reed, C. M.", "Regimbau, T.", "Rei, L.", "Reid, S.", "Reitze, D. H.", "Rew, H.", "Reyes, S. D.", "Ricci, F.", "Riles, K.", "Rizzo, M.", "Robertson, N. A.", "Robie, R.", "Robinet, F.", "Rocchi, A.", "Rolland, L.", "Rollins, J. G.", "Roma, V. J.", "Romano, R.", "Romanov, G.", "Romie, J. H.", "Rosińska, D.", "Rowan, S.", "Rüdiger, A.", "Ruggi, P.", "Ryan, K.", "Sachdev, S.", "Sadecki, T.", "Sadeghian, L.", "Sakellariadou, M.", "Salconi, L.", "Saleem, M.", "Salemi, F.", "Samajdar, A.", "Sammut, L.", "Sanchez, E. J.", "Sandberg, V.", "Sandeen, B.", "Sanders, J. R.", "Sassolas, B.", "Sathyaprakash, B. S.", "Saulson, P. R.", "Sauter, O. E. S.", "Savage, R. L.", "Sawadsky, A.", "Schale, P.", "Schilling, R.", "Schmidt, J.", "Schmidt, P.", "Schnabel, R.", "Schofield, R. M. S.", "Schönbeck, A.", "Schreiber, E.", "Schuette, D.", "Schutz, B. F.", "Scott, J.", "Scott, S. M.", "Sellers, D.", "Sengupta, A. S.", "Sentenac, D.", "Sequino, V.", "Sergeev, A.", "Setyawati, Y.", "Shaddock, D. A.", "Shaffer, T.", "Shahriar, M. S.", "Shaltev, M.", "Shapiro, B.", "Shawhan, P.", "Sheperd, A.", "Shoemaker, D. H.", "Shoemaker, D. M.", "Siellez, K.", "Siemens, X.", "Sieniawska, M.", "Sigg, D.", "Silva, A. D.", "Singer, A.", "Singer, L. P.", "Singh, A.", "Singh, R.", "Singhal, A.", "Sintes, A. M.", "Slagmolen, B. J. J.", "Smith, J. R.", "Smith, N. D.", "Smith, R. J. E.", "Son, E. J.", "Sorazu, B.", "Sorrentino, F.", "Souradeep, T.", "Srivastava, A. K.", "Staley, A.", "Steinke, M.", "Steinlechner, J.", "Steinlechner, S.", "Steinmeyer, D.", "Stephens, B. C.", "Stone, R.", "Strain, K. A.", "Straniero, N.", "Stratta, G.", "Strauss, N. A.", "Strigin, S.", "Sturani, R.", "Stuver, A. L.", "Summerscales, T. Z.", "Sun, L.", "Sunil, S.", "Sutton, P. J.", "Swinkels, B. L.", "Szczepańczyk, M. J.", "Tacca, M.", "Talukder, D.", "Tanner, D. B.", "Tápai, M.", "Tarabrin, S. P.", "Taracchini, A.", "Taylor, R.", "Theeg, T.", "Thirugnanasambandam, M. P.", "Thomas, E. G.", "Thomas, M.", "Thomas, P.", "Thorne, K. A.", "Thrane, E.", "Tiwari, S.", "Tiwari, V.", "Tokmakov, K. V.", "Toland, K.", "Tomlinson, C.", "Tonelli, M.", "Tornasi, Z.", "Torres, C. V.", "Torrie, C. I.", "Töyrä, D.", "Travasso, F.", "Traylor, G.", "Trifirò, D.", "Tringali, M. C.", "Trozzo, L.", "Tse, M.", "Turconi, M.", "Tuyenbayev, D.", "Ugolini, D.", "Unnikrishnan, C. S.", "Urban, A. L.", "Usman, S. A.", "Vahlbruch, H.", "Vajente, G.", "Valdes, G.", "van Bakel, N.", "van Beuzekom, M.", "van den Brand, J. F. J.", "Van Den Broeck, C.", "Vander-Hyde, D. C.", "van der Schaaf, L.", "van Heijningen, J. V.", "van Veggel, A. A.", "Vardaro, M.", "Vass, S.", "Vasúth, M.", "Vaulin, R.", "Vecchio, A.", "Vedovato, G.", "Veitch, J.", "Veitch, P. J.", "Venkateswara, K.", "Verkindt, D.", "Vetrano, F.", "Viceré, A.", "Vinciguerra, S.", "Vine, D. J.", "Vinet, J. -Y.", "Vitale, S.", "Vo, T.", "Vocca, H.", "Vorvick, C.", "Voss, D. V.", "Vousden, W. D.", "Vyatchanin, S. P.", "Wade, A. R.", "Wade, L. E.", "Wade, M.", "Walker, M.", "Wallace, L.", "Walsh, S.", "Wang, G.", "Wang, H.", "Wang, M.", "Wang, X.", "Wang, Y.", "Ward, R. L.", "Warner, J.", "Was, M.", "Weaver, B.", "Wei, L. -W.", "Weinert, M.", "Weinstein, A. J.", "Weiss, R.", "Wen, L.", "Weßels, P.", "Westphal, T.", "Wette, K.", "Whelan, J. T.", "Whiting, B. F.", "Williams, R. D.", "Williamson, A. R.", "Willis, J. L.", "Willke, B.", "Wimmer, M. H.", "Winkler, W.", "Wipf, C. C.", "Wittel, H.", "Woan, G.", "Woehler, J.", "Worden, J.", "Wright, J. L.", "Wu, D. S.", "Wu, G.", "Yablon, J.", "Yam, W.", "Yamamoto, H.", "Yancey, C. C.", "Yu, H.", "Yvert, M.", "Zadrożny, A.", "Zangrando, L.", "Zanolin, M.", "Zendri, J. -P.", "Zevin, M.", "Zhang, L.", "Zhang, M.", "Zhang, Y.", "Zhao, C.", "Zhou, M.", "Zhou, Z.", "Zhu, X. J.", "Zucker, M. E.", "Zuraw, S. E.", "Zweizig, J.", "LIGO Scientific Collaboration", "Virgo Collaboration" ]
2016ApJ...832L..21A
[ "Upper Limits on the Rates of Binary Neutron Star and Neutron Star-Black Hole Mergers from Advanced LIGO’s First Observing Run" ]
165
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Angelo, I-80126 Napoli, Italy", "University of Florida, Gainesville, FL 32611, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "University of Sannio at Benevento, I-82100 Benevento, Italy ; INFN, Sezione di Napoli, I-80100 Napoli, Italy", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Instituto Nacional de Pesquisas Espaciais, 12227-010 São José dos Campos, São Paulo, Brazil", "INFN, Gran Sasso Science Institute, I-67100 L'Aquila, Italy ; INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy", "Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India", "International Centre for Theoretical Sciences, Tata Institute of Fundamental Research, Bangalore 560012, India", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA ; Leibniz Universität Hannover, D-30167 Hannover, Germany", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Australian National University, Canberra, ACT 0200, Australia", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "The University of Mississippi, University, MS 38677, USA", "California State University Fullerton, Fullerton, CA 92831, USA", "LAL, Univ. 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Angelo, I-80126 Napoli, Italy", "Laboratoire Kastler Brossel, UPMC-Sorbonne Universités, CNRS, ENS-PSL Research University, Collège de France, F-75005 Paris, France", "University of Birmingham, Birmingham B15 2TT, UK", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "INFN, Sezione di Napoli, Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy ; Università di Napoli “Federico II,” Complesso Universitario di Monte S. 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Angelo, I-80126 Napoli, Italy", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "INFN, Sezione di Roma, I-00185 Roma, Italy; Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "The University of Mississippi, University, MS 38677, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "LIGO Hanford Observatory, Richland, WA 99352, USA", "The University of Sheffield, Sheffield S10 2TN, UK", "Carleton College, Northfield, MN 55057, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA; University of Florida, Gainesville, FL 32611, USA", "University of Florida, Gainesville, FL 32611, USA", "Caltech CaRT, Pasadena, CA 91125, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "The Pennsylvania State University, University Park, PA 16802, USA", "Columbia University, New York, NY 10027, USA", "INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy; Università di Roma Tor Vergata, I-00133 Roma, Italy", "Syracuse University, Syracuse, NY 13244, USA", "Cardiff University, Cardiff CF24 3AA, UK", "Tsinghua University, Beijing 100084, China", "University of Western Australia, Crawley, WA 6009, Australia", "INFN, Sezione di Genova, I-16146 Genova, Italy", "University of Chicago, Chicago, IL 60637, USA", "University of Birmingham, Birmingham B15 2TT, UK", "Montclair State University, Montclair, NJ 07043, USA", "Cardiff University, Cardiff CF24 3AA, UK", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Stanford University, Stanford, CA 94305, USA", "MTA Eötvös University, “Lendulet” Astrophysics Research Group, Budapest 1117, Hungary", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Instituto Nacional de Pesquisas Espaciais, 12227-010 São José dos Campos, São Paulo, Brazil", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "APC, AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/Irfu, Observatoire de Paris, Sorbonne Paris Cité, F-75205 Paris Cedex 13, France", "Syracuse University, Syracuse, NY 13244, USA", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France; National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "INFN, Sezione di Pisa, I-56127 Pisa, Italy", "MTA Eötvös University, “Lendulet” Astrophysics Research Group, Budapest 1117, Hungary", "University of Birmingham, Birmingham B15 2TT, UK", "University of Oregon, Eugene, OR 97403, USA", "LAL, Univ. Paris-Sud, CNRS/IN2P3, Université Paris-Saclay, Orsay, France", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "University of Florida, Gainesville, FL 32611, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "The University of Mississippi, University, MS 38677, USA", "School of Mathematics, University of Edinburgh, Edinburgh EH9 3FD, UK", "Università di Perugia, I-06123 Perugia, Italy", "Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India", "INFN, Sezione di Napoli, Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy; Università di Napoli “Federico II,” Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy", "Institute for Plasma Research, Bhat, Gandhinagar 382428, India; Indian Institute of Technology, Gandhinagar Ahmedabad Gujarat 382424, India", "NASA/Goddard Space Flight Center, Greenbelt, MD 20771, USA", "INFN, Sezione di Genova, I-16146 Genova, Italy", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "INFN, Sezione di Pisa, I-56127 Pisa, Italy", "RRCAT, Indore MP 452013, India", "University of Szeged, Dóm tér 9, Szeged 6720, Hungary", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "International Centre for Theoretical Sciences, Tata Institute of Fundamental Research, Bangalore 560012, India", "International Centre for Theoretical Sciences, Tata Institute of Fundamental Research, Bangalore 560012, India", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands; Department of Astrophysics/IMAPP, Radboud University Nijmegen, P.O. Box 9010, 6500 GL Nijmegen, The Netherlands", "Louisiana State University, Baton Rouge, LA 70803, USA; LIGO Livingston Observatory, Livingston, LA 70754, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Embry-Riddle Aeronautical University, Prescott, AZ 86301, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "LIGO Hanford Observatory, Richland, WA 99352, USA", "University of Florida, Gainesville, FL 32611, USA", "MTA Eötvös University, “Lendulet” Astrophysics Research Group, Budapest 1117, Hungary", "Louisiana State University, Baton Rouge, LA 70803, USA", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Tata Institute of Fundamental Research, Mumbai 400005, India", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Faculty of Physics, Lomonosov Moscow State University, Moscow 119991, Russia", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "INFN, Sezione di Napoli, Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy; INAF, Osservatorio Astronomico di Capodimonte, I-80131 Napoli, Italy", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "University of Maryland, College Park, MD 20742, USA", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "University of Birmingham, Birmingham B15 2TT, UK", "Department of Astrophysics/IMAPP, Radboud University Nijmegen, P.O. Box 9010, 6500 GL Nijmegen, The Netherlands", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "Tsinghua University, Beijing 100084, China", "Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India", "Institute for Plasma Research, Bhat, Gandhinagar 382428, India", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "University of Michigan, Ann Arbor, MI 48109, USA", "California State University Fullerton, Fullerton, CA 92831, USA", "Washington State University, Pullman, WA 99164, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Tata Institute of Fundamental Research, Mumbai 400005, India", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO Hanford Observatory, Richland, WA 99352, USA", "The Pennsylvania State University, University Park, PA 16802, USA", "Cardiff University, Cardiff CF24 3AA, UK", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Louisiana State University, Baton Rouge, LA 70803, USA", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "American University, Washington, D.C. 20016, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "University of Florida, Gainesville, FL 32611, USA", "University of Birmingham, Birmingham B15 2TT, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Laboratoire Kastler Brossel, UPMC-Sorbonne Universités, CNRS, ENS-PSL Research University, Collège de France, F-75005 Paris, France", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "LAL, Univ. Paris-Sud, CNRS/IN2P3, Université Paris-Saclay, Orsay, France", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Rochester Institute of Technology, Rochester, NY 14623, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "University of Chicago, Chicago, IL 60637, USA", "Cardiff University, Cardiff CF24 3AA, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "University of Western Australia, Crawley, WA 6009, Australia", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "National Tsing Hua University, Hsinchu City, 30013 Taiwan", "NCSA, University of Illinois at Urbana-Champaign, Urbana, IL 61801, USA", "LAL, Univ. Paris-Sud, CNRS/IN2P3, Université Paris-Saclay, Orsay, France", "Embry-Riddle Aeronautical University, Prescott, AZ 86301, USA", "Universitat de les Illes Balears, IAC3—IEEC, E-07122 Palma de Mallorca, Spain", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Texas Tech University, Lubbock, TX 79409, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Laboratoire Kastler Brossel, UPMC-Sorbonne Universités, CNRS, ENS-PSL Research University, Collège de France, F-75005 Paris, France", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "International Centre for Theoretical Sciences, Tata Institute of Fundamental Research, Bangalore 560012, India", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Laboratoire Kastler Brossel, UPMC-Sorbonne Universités, CNRS, ENS-PSL Research University, Collège de France, F-75005 Paris, France", "Korea Institute of Science and Technology Information, Daejeon 305-806, Korea", "Center for Relativistic Astrophysics and School of Physics, Georgia Institute of Technology, Atlanta, GA 30332, USA", "University of Białystok, 15-424 Białystok, Poland", "SUPA, University of Strathclyde, Glasgow G1 1XQ, UK", "University of Western Australia, Crawley, WA 6009, Australia", "Universitat de les Illes Balears, IAC3—IEEC, E-07122 Palma de Mallorca, Spain", "Louisiana State University, Baton Rouge, LA 70803, USA", "University of Southampton, Southampton SO17 1BJ, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "University of Western Australia, Crawley, WA 6009, Australia", "IISER-TVM, CET Campus, Trivandrum Kerala 695016, India", "Cardiff University, Cardiff CF24 3AA, UK", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "The University of Mississippi, University, MS 38677, USA", "Korea Institute of Science and Technology Information, Daejeon 305-806, Korea", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "University of Oregon, Eugene, OR 97403, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Louisiana State University, Baton Rouge, LA 70803, USA; European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Leibniz Universität Hannover, D-30167 Hannover, Germany", "University of Western Australia, Crawley, WA 6009, Australia", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "Canadian Institute for Theoretical Astrophysics, University of Toronto, Toronto, ON M5S 3H8, Canada", "Universitat de les Illes Balears, IAC3—IEEC, E-07122 Palma de Mallorca, Spain", "Syracuse University, Syracuse, NY 13244, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "The University of Sheffield, Sheffield S10 2TN, UK", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "Faculty of Physics, Lomonosov Moscow State University, Moscow 119991, Russia", "INFN, Gran Sasso Science Institute, I-67100 L'Aquila, Italy", "Cardiff University, Cardiff CF24 3AA, UK", "Institute for Plasma Research, Bhat, Gandhinagar 382428, India", "Institute of Applied Physics, Nizhny Novgorod, 603950, Russia", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Korea Institute of Science and Technology Information, Daejeon 305-806, Korea", "Korea Institute of Science and Technology Information, Daejeon 305-806, Korea", "Pusan National University, Busan 609-735, Korea", "Hanyang University, Seoul 133-791, Korea", "Stanford University, Stanford, CA 94305, USA", "University of Adelaide, Adelaide, SA 5005, Australia", "Pusan National University, Busan 609-735, Korea", "Center for Relativistic Astrophysics and School of Physics, Georgia Institute of Technology, Atlanta, GA 30332, USA", "University of Adelaide, Adelaide, SA 5005, Australia", "LIGO Hanford Observatory, Richland, WA 99352, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "Universität Hamburg, D-22761 Hamburg, Germany", "University of Florida, Gainesville, FL 32611, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Universität Hamburg, D-22761 Hamburg, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Astronomical Observatory Warsaw University, 00-478 Warsaw, Poland", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "NCBJ, 05-400 Świerk-Otwock, Poland ; IM-PAN, 00-956 Warsaw, Poland", "Leibniz Universität Hannover, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Canadian Institute for Theoretical Astrophysics, University of Toronto, Toronto, ON M5S 3H8, Canada", "Institute for Plasma Research, Bhat, Gandhinagar 382428, India", "National Tsing Hua University, Hsinchu City, 30013 Taiwan", "NCBJ, 05-400 Świerk-Otwock, Poland", "Syracuse University, Syracuse, NY 13244, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Rochester Institute of Technology, Rochester, NY 14623, USA", "Stanford University, Stanford, CA 94305, USA", "Monash University, Melbourne, VIC 3800, Australia", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "INFN, Sezione di Padova, I-35131 Padova, Italy", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "APC, AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/Irfu, Observatoire de Paris, Sorbonne Paris Cité, F-75205 Paris Cedex 13, France; Tsinghua University, Beijing 100084, China", "Pusan National University, Busan 609-735, Korea", "Hanyang University, Seoul 133-791, Korea", "Seoul National University, Seoul 151-742, Korea", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Syracuse University, Syracuse, NY 13244, USA", "Università di Trento, Dipartimento di Fisica, I-38123 Povo, Trento, Italy ; INFN, Trento Institute for Fundamental Physics and Applications, I-38123 Povo, Trento, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LAL, Univ. Paris-Sud, CNRS/IN2P3, Université Paris-Saclay, Orsay, France", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "Monash University, Melbourne, VIC 3800, Australia", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "The Chinese University of Hong Kong, Shatin, NT, Hong Kong", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "University of Alabama in Huntsville, Huntsville, AL 35899, USA", "SUPA, University of Strathclyde, Glasgow G1 1XQ, UK", "University of Massachusetts-Amherst, Amherst, MA 01003, USA", "Cardiff University, Cardiff CF24 3AA, UK", "Syracuse University, Syracuse, NY 13244, USA", "INFN, Gran Sasso Science Institute, I-67100 L'Aquila, Italy; INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy", "ESPCI, CNRS, F-75005 Paris, France", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "University of Western Australia, Crawley, WA 6009, Australia", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Louisiana State University, Baton Rouge, LA 70803, USA", "Syracuse University, Syracuse, NY 13244, USA", "Syracuse University, Syracuse, NY 13244, USA", "Washington State University, Pullman, WA 99164, USA", "INFN, Sezione di Roma, I-00185 Roma, Italy", "ESPCI, CNRS, F-75005 Paris, France", "INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy; Università di Roma Tor Vergata, I-00133 Roma, Italy", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "University of Minnesota, Minneapolis, MN 55455, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Australian National University, Canberra, ACT 0200, Australia", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "INFN, Sezione di Perugia, I-06123 Perugia, Italy ; Università di Camerino, Dipartimento di Fisica, I-62032 Camerino, Italy", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "Columbia University, New York, NY 10027, USA", "Columbia University, New York, NY 10027, USA", "Stanford University, Stanford, CA 94305, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "Syracuse University, Syracuse, NY 13244, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Columbia University, New York, NY 10027, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Washington State University, Pullman, WA 99164, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Australian National University, Canberra, ACT 0200, Australia", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Southern University and A&amp;M College, Baton Rouge, LA 70813, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Australian National University, Canberra, ACT 0200, Australia", "Australian National University, Canberra, ACT 0200, Australia", "West Virginia University, Morgantown, WV 26506, USA", "The Pennsylvania State University, University Park, PA 16802, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "The University of Melbourne, Parkville, Victoria 3010, Australia", "LIGO Hanford Observatory, Richland, WA 99352, USA", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "The Pennsylvania State University, University Park, PA 16802, USA", "Laboratoire Kastler Brossel, UPMC-Sorbonne Universités, CNRS, ENS-PSL Research University, Collège de France, F-75005 Paris, France", "University of Minnesota, Minneapolis, MN 55455, USA", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "University of Birmingham, Birmingham B15 2TT, UK", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France", "University of Birmingham, Birmingham B15 2TT, UK", "College of William and Mary, Williamsburg, VA 23187, USA", "INFN, Sezione di Napoli, Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy; Università di Napoli “Federico II,” Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy", "University of Florida, Gainesville, FL 32611, USA; INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Montana State University, Bozeman, MT 59717, USA", "INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "Instituto de Física Teórica, University Estadual Paulista/ICTP South American Institute for Fundamental Research, São Paulo SP 01140-070, Brazil", "International Centre for Theoretical Sciences, Tata Institute of Fundamental Research, Bangalore 560012, India", "Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India", "Faculty of Physics, Lomonosov Moscow State University, Moscow 119991, Russia", "University of Florida, Gainesville, FL 32611, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "INFN, Sezione di Pisa, I-56127 Pisa, Italy", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "Montclair State University, Montclair, NJ 07043, USA", "University of Cambridge, Cambridge CB2 1TN, UK", "LIGO Hanford Observatory, Richland, WA 99352, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "University of Birmingham, Birmingham B15 2TT, UK", "University of Florida, Gainesville, FL 32611, USA", "Cardiff University, Cardiff CF24 3AA, UK", "International Centre for Theoretical Sciences, Tata Institute of Fundamental Research, Bangalore 560012, India", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "University of Adelaide, Adelaide, SA 5005, Australia", "Columbia University, New York, NY 10027, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "University of Florida, Gainesville, FL 32611, USA", "INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy; Università di Roma Tor Vergata, I-00133 Roma, Italy", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "IISER-Kolkata, Mohanpur, West Bengal 741252, India", "University of Massachusetts-Amherst, Amherst, MA 01003, USA", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands; Department of Astrophysics/IMAPP, Radboud University Nijmegen, P.O. Box 9010, 6500 GL Nijmegen, The Netherlands", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Università degli Studi di Genova, I-16146 Genova, Italy ; INFN, Sezione di Genova, I-16146 Genova, Italy", "University of Michigan, Ann Arbor, MI 48109, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Australian National University, Canberra, ACT 0200, Australia", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands; Department of Astrophysics/IMAPP, Radboud University Nijmegen, P.O. Box 9010, 6500 GL Nijmegen, The Netherlands", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "Syracuse University, Syracuse, NY 13244, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "Rutherford Appleton Laboratory, HSIC, Chilton, Didcot, Oxon OX11 0QX, UK", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Whitman College, 345 Boyer Avenue, Walla Walla, WA 99362, USA", "National Institute for Mathematical Sciences, Daejeon 305-390, Korea", "National Institute for Mathematical Sciences, Daejeon 305-390, Korea", "Cardiff University, Cardiff CF24 3AA, UK", "Universitat de les Illes Balears, IAC3—IEEC, E-07122 Palma de Mallorca, Spain", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Rochester Institute of Technology, Rochester, NY 14623, USA", "University of Adelaide, Adelaide, SA 5005, Australia", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Texas Tech University, Lubbock, TX 79409, USA", "IISER-TVM, CET Campus, Trivandrum Kerala 695016, India", "RRCAT, Indore MP 452013, India", "University of Oregon, Eugene, OR 97403, USA", "Institute of Applied Physics, Nizhny Novgorod, 603950, Russia", "INFN, Sezione di Roma, I-00185 Roma, Italy", "Universität Hamburg, D-22761 Hamburg, Germany", "National Tsing Hua University, Hsinchu City, 30013 Taiwan", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "Cardiff University, Cardiff CF24 3AA, UK", "RRCAT, Indore MP 452013, India", "INFN, Sezione di Pisa, I-56127 Pisa, Italy; European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA; Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "Stanford University, Stanford, CA 94305, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Stanford University, Stanford, CA 94305, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France; Université de Lyon, F-69361 Lyon, France", "Syracuse University, Syracuse, NY 13244, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Hobart and William Smith Colleges, Geneva, NY 14456, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "University of Sannio at Benevento, I-82100 Benevento, Italy ; INFN, Sezione di Napoli, I-80100 Napoli, Italy", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France", "University of Sannio at Benevento, I-82100 Benevento, Italy ; INFN, Sezione di Napoli, I-80100 Napoli, Italy", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India", "Cardiff University, Cardiff CF24 3AA, UK", "University of Minnesota, Minneapolis, MN 55455, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "University of Sannio at Benevento, I-82100 Benevento, Italy ; INFN, Sezione di Napoli, I-80100 Napoli, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Università di Trento, Dipartimento di Fisica, I-38123 Povo, Trento, Italy ; INFN, Trento Institute for Fundamental Physics and Applications, I-38123 Povo, Trento, Italy", "Faculty of Physics, Lomonosov Moscow State University, Moscow 119991, Russia", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "INFN, Sezione di Perugia, I-06123 Perugia, Italy", "INFN, Sezione di Roma, I-00185 Roma, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "University of Western Australia, Crawley, WA 6009, Australia", "Monash University, Melbourne, VIC 3800, Australia", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "University of Oregon, Eugene, OR 97403, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Australian National University, Canberra, ACT 0200, Australia", "LIGO Hanford Observatory, Richland, WA 99352, USA", "MTA Eötvös University, “Lendulet” Astrophysics Research Group, Budapest 1117, Hungary", "RRCAT, Indore MP 452013, India", "RRCAT, Indore MP 452013, India", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "Università di Roma Tor Vergata, I-00133 Roma, Italy", "California State University Fullerton, Fullerton, CA 92831, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "INFN, Sezione di Genova, I-16146 Genova, Italy", "SUPA, University of the West of Scotland, Paisley PA1 2BE, UK", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA; University of Florida, Gainesville, FL 32611, USA", "College of William and Mary, Williamsburg, VA 23187, USA", "Syracuse University, Syracuse, NY 13244, USA", "INFN, Sezione di Roma, I-00185 Roma, Italy; Università di Roma “La Sapienza,” I-00185 Roma, Italy", "University of Michigan, Ann Arbor, MI 48109, USA", "Rochester Institute of Technology, Rochester, NY 14623, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA; SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "LAL, Univ. Paris-Sud, CNRS/IN2P3, Université Paris-Saclay, Orsay, France", "INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "University of Oregon, Eugene, OR 97403, USA", "Università di Salerno, Fisciano, I-84084 Salerno, Italy ; INFN, Sezione di Napoli, Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy", "College of William and Mary, Williamsburg, VA 23187, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "CAMK-PAN, 00-716 Warsaw, Poland; Janusz Gil Institute of Astronomy, University of Zielona Góra, 65-265 Zielona Góra, Poland", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "LIGO Hanford Observatory, Richland, WA 99352, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "King's College London, University of London, London WC2R 2LS, UK", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "IISER-TVM, CET Campus, Trivandrum Kerala 695016, India", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "IISER-Kolkata, Mohanpur, West Bengal 741252, India", "Monash University, Melbourne, VIC 3800, Australia", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "Syracuse University, Syracuse, NY 13244, USA", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France", "Cardiff University, Cardiff CF24 3AA, UK", "Syracuse University, Syracuse, NY 13244, USA", "University of Michigan, Ann Arbor, MI 48109, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Leibniz Universität Hannover, D-30167 Hannover, Germany", "University of Oregon, Eugene, OR 97403, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA; Caltech CaRT, Pasadena, CA 91125, USA", "Universität Hamburg, D-22761 Hamburg, Germany", "University of Oregon, Eugene, OR 97403, USA", "Universität Hamburg, D-22761 Hamburg, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany; Cardiff University, Cardiff CF24 3AA, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Australian National University, Canberra, ACT 0200, Australia", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Indian Institute of Technology, Gandhinagar Ahmedabad Gujarat 382424, India", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy; Università di Roma Tor Vergata, I-00133 Roma, Italy", "Institute of Applied Physics, Nizhny Novgorod, 603950, Russia", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands; Department of Astrophysics/IMAPP, Radboud University Nijmegen, P.O. Box 9010, 6500 GL Nijmegen, The Netherlands", "Australian National University, Canberra, ACT 0200, Australia", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Stanford University, Stanford, CA 94305, USA", "University of Maryland, College Park, MD 20742, USA", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Center for Relativistic Astrophysics and School of Physics, Georgia Institute of Technology, Atlanta, GA 30332, USA", "Center for Relativistic Astrophysics and School of Physics, Georgia Institute of Technology, Atlanta, GA 30332, USA", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "CAMK-PAN, 00-716 Warsaw, Poland", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Instituto Nacional de Pesquisas Espaciais, 12227-010 São José dos Campos, São Paulo, Brazil", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "NASA/Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany; Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "Louisiana State University, Baton Rouge, LA 70803, USA", "INFN, Gran Sasso Science Institute, I-67100 L'Aquila, Italy", "Universitat de les Illes Balears, IAC3—IEEC, E-07122 Palma de Mallorca, Spain", "Australian National University, Canberra, ACT 0200, Australia", "California State University Fullerton, Fullerton, CA 92831, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "National Institute for Mathematical Sciences, Daejeon 305-390, Korea", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "INFN, Sezione di Genova, I-16146 Genova, Italy", "Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India", "Institute for Plasma Research, Bhat, Gandhinagar 382428, India", "Columbia University, New York, NY 10027, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Laboratoire des Matériaux Avancés (LMA), CNRS/IN2P3, F-69622 Villeurbanne, France", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "Carleton College, Northfield, MN 55057, USA", "Faculty of Physics, Lomonosov Moscow State University, Moscow 119991, Russia", "Instituto de Física Teórica, University Estadual Paulista/ICTP South American Institute for Fundamental Research, São Paulo SP 01140-070, Brazil", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Andrews University, Berrien Springs, MI 49104, USA", "The University of Melbourne, Parkville, Victoria 3010, Australia", "Institute for Plasma Research, Bhat, Gandhinagar 382428, India", "Cardiff University, Cardiff CF24 3AA, UK", "European Gravitational Observatory (EGO), I-56021 Cascina, Pisa, Italy", "Embry-Riddle Aeronautical University, Prescott, AZ 86301, USA", "APC, AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/Irfu, Observatoire de Paris, Sorbonne Paris Cité, F-75205 Paris Cedex 13, France", "University of Oregon, Eugene, OR 97403, USA", "University of Florida, Gainesville, FL 32611, USA", "University of Szeged, Dóm tér 9, Szeged 6720, Hungary", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "University of Birmingham, Birmingham B15 2TT, UK", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "LIGO Hanford Observatory, Richland, WA 99352, USA", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Monash University, Melbourne, VIC 3800, Australia", "INFN, Gran Sasso Science Institute, I-67100 L'Aquila, Italy; INFN, Trento Institute for Fundamental Physics and Applications, I-38123 Povo, Trento, Italy", "Cardiff University, Cardiff CF24 3AA, UK", "SUPA, University of Strathclyde, Glasgow G1 1XQ, UK", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "The University of Sheffield, Sheffield S10 2TN, UK", "Università di Pisa, I-56127 Pisa, Italy ; INFN, Sezione di Pisa, I-56127 Pisa, Italy", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "University of Birmingham, Birmingham B15 2TT, UK", "Università di Perugia, I-06123 Perugia, Italy; INFN, Sezione di Perugia, I-06123 Perugia, Italy", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "The University of Mississippi, University, MS 38677, USA", "Università di Trento, Dipartimento di Fisica, I-38123 Povo, Trento, Italy ; INFN, Trento Institute for Fundamental Physics and Applications, I-38123 Povo, Trento, Italy", "INFN, Sezione di Pisa, I-56127 Pisa, Italy; Università di Siena, I-53100 Siena, Italy", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "Trinity University, San Antonio, TX 78212, USA", "Tata Institute of Fundamental Research, Mumbai 400005, India", "University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA", "Syracuse University, Syracuse, NY 13244, USA", "Leibniz Universität Hannover, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "The University of Texas Rio Grande Valley, Brownsville, TX 78520, USA", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands; VU University Amsterdam, 1081 HV Amsterdam, The Netherlands", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "Syracuse University, Syracuse, NY 13244, USA", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Università di Padova, Dipartimento di Fisica e Astronomia, I-35131 Padova, Italy ; INFN, Sezione di Padova, I-35131 Padova, Italy", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Wigner RCP, RMKI, H-1121 Budapest, Konkoly Thege Miklós út 29-33, Hungary", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "University of Birmingham, Birmingham B15 2TT, UK", "INFN, Sezione di Padova, I-35131 Padova, Italy", "University of Birmingham, Birmingham B15 2TT, UK", "University of Adelaide, Adelaide, SA 5005, Australia", "University of Washington, Seattle, WA 98195, USA", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "Università degli Studi di Urbino “Carlo Bo,” I-61029 Urbino, Italy ; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "University of Birmingham, Birmingham B15 2TT, UK", "SUPA, University of the West of Scotland, Paisley PA1 2BE, UK", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Syracuse University, Syracuse, NY 13244, USA", "Università di Perugia, I-06123 Perugia, Italy; INFN, Sezione di Perugia, I-06123 Perugia, Italy", "LIGO Hanford Observatory, Richland, WA 99352, USA", "University of Florida, Gainesville, FL 32611, USA", "University of Birmingham, Birmingham B15 2TT, UK", "Faculty of Physics, Lomonosov Moscow State University, Moscow 119991, Russia", "Australian National University, Canberra, ACT 0200, Australia", "Kenyon College, Gambier, OH 43022, USA", "Kenyon College, Gambier, OH 43022, USA", "Louisiana State University, Baton Rouge, LA 70803, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-14476 Potsdam-Golm, Germany", "INFN, Gran Sasso Science Institute, I-67100 L'Aquila, Italy; INFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy", "University of Birmingham, Birmingham B15 2TT, UK", "University of Birmingham, Birmingham B15 2TT, UK", "Tsinghua University, Beijing 100084, China", "University of Western Australia, Crawley, WA 6009, Australia", "Australian National University, Canberra, ACT 0200, Australia", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "LIGO Hanford Observatory, Richland, WA 99352, USA", "Artemis, Université Côte d'Azur, CNRS, Observatoire Côte d'Azur, CS 34229, Nice cedex 4, France", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "University of Western Australia, Crawley, WA 6009, Australia", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "Rochester Institute of Technology, Rochester, NY 14623, USA", "University of Florida, Gainesville, FL 32611, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Cardiff University, Cardiff CF24 3AA, UK", "Abilene Christian University, Abilene, TX 79699, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany; Leibniz Universität Hannover, D-30167 Hannover, Germany", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO Hanford Observatory, Richland, WA 99352, USA", "SUPA, University of Glasgow, Glasgow G12 8QQ, UK", "Albert-Einstein-Institut, Max-Planck-Institut für Gravitationsphysik, D-30167 Hannover, Germany", "LIGO Livingston Observatory, Livingston, LA 70754, USA", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "University of Maryland, College Park, MD 20742, USA", "LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Laboratoire d'Annecy-le-Vieux de Physique des Particules (LAPP), Université Savoie Mont Blanc, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France", "NCBJ, 05-400 Świerk-Otwock, Poland", "INFN, Sezione di Padova, I-35131 Padova, Italy", "Embry-Riddle Aeronautical University, Prescott, AZ 86301, USA", "INFN, Sezione di Padova, I-35131 Padova, Italy", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "College of William and Mary, Williamsburg, VA 23187, USA", "Rochester Institute of Technology, Rochester, NY 14623, USA", "University of Western Australia, Crawley, WA 6009, Australia", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "Center for Interdisciplinary Exploration &amp; Research in Astrophysics (CIERA), Northwestern University, Evanston, IL 60208, USA", "University of Western Australia, Crawley, WA 6009, Australia", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA; LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "University of Massachusetts-Amherst, Amherst, MA 01003, USA", "LIGO, California Institute of Technology, Pasadena, CA 91125, USA", "-", "-" ]
[ "2015mbhe.confE..53F", "2016ApJ...829L..15S", "2016CQGra..33u5004U", "2016PhRvD..94e3013L", "2016PhRvX...6d1015A", "2016arXiv161201471C", "2017A&A...604A..55M", "2017AcA....67...37C", "2017ApJ...836..230C", "2017ApJ...837...57D", "2017ApJ...837...67C", "2017ApJ...841...89A", "2017ApJ...844L..22L", "2017ApJ...846...62S", "2017ApJ...846..142K", "2017ApJ...847...47A", "2017ApJ...848L..17C", "2017ApJ...848L..26S", "2017ApJ...848L..34M", "2017ApJ...849..114Z", "2017ApJ...849..118N", "2017ApJ...849..153F", "2017ApJ...849L..14G", "2017ApJ...850L..24G", "2017ApJ...851L..48Y", "2017AstL...43..516B", "2017CQGra..34j4001R", "2017CQGra..34j5014D", "2017CQGra..34l4001L", "2017CQGra..34x5003C", "2017JPhCS.837a2010P", "2017MNRAS.464.2622Y", "2017MNRAS.466.2085M", "2017MNRAS.470..350Y", "2017MNRAS.471.4488S", "2017NatAs...1E.112K", "2017Natur.551...75S", "2017Natur.551...80K", "2017PhRvD..95d4039K", "2017PhRvD..95f2002B", "2017PhRvD..95f3016C", "2017PhRvD..95f4013S", "2017PhRvD..95j3007W", "2017PhRvD..95l4006D", "2017PhRvD..95l4008Z", "2017PhRvD..95l4042H", "2017PhRvD..96b2001A", "2017PhRvD..96d3001T", "2017PhRvD..96d3019K", "2017PhRvD..96h4019S", "2017PhRvD..96h4063R", "2017PhRvD..96j4015K", "2017PhRvD..96l1501D", "2017PhRvD..96l4005B", "2017PhRvL.119p1101A", "2017Sci...358.1559K", "2017Sci...358.1583K", "2017arXiv170103952H", "2017arXiv170501845D", "2017arXiv171108743R", "2018ApJ...852L...3S", "2018ApJ...853..123S", "2018ApJ...854L..25P", "2018ApJ...855...34G", "2018ApJ...855...99C", "2018ApJ...857..128J", "2018ApJ...858...79W", "2018ApJ...860....5G", "2018ApJ...861...85A", "2018ApJ...864...22A", "2018ApJ...867...39K", "2018ApJS..234...34P", "2018ApJS..234...41W", "2018CQGra..35f5009A", "2018JKPS...72....1C", "2018LRR....21....3A", "2018MNRAS.473.1186E", "2018MNRAS.475.1331T", "2018MNRAS.477..639B", "2018MNRAS.478.3298W", "2018MNRAS.479..601D", "2018MNRAS.479.4391M", "2018MNRAS.480.2011G", "2018MNRAS.481.1597G", "2018PASJ...70...39Y", "2018PhRvD..97d3015P", "2018PhRvD..97e5016B", "2018PhRvD..97j4064M", "2018PhRvD..98b4050N", "2018PhRvD..98d3002Z", "2018PhRvD..98f3501J", "2018PhRvD..98h3007N", "2018PhRvD..98j3024M", "2018PhRvD..98j4055N", "2018PhRvD..98l3021H", "2018PhRvD..98l4014J", "2018PhRvD..98l4030S", "2018PhRvL.120c1102B", "2018PhRvL.120i1101A", "2018PhRvL.121b1303V", "2018arXiv180307965D", "2018arXiv181001421G", "2018arXiv181205225C", "2019AnP...53100365A", "2019ApJ...871...90B", "2019ApJ...871..178G", "2019ApJ...872..195N", "2019ApJ...878L..17M", "2019ApJ...880...55M", "2019ApJ...881..157B", "2019BAAS...51c.449W", "2019CQGra..36n3001B", "2019IAUS..346..433C", "2019ICRC...36..803T", "2019MNRAS.482.5012C", "2019MNRAS.484.1506L", "2019MNRAS.484.4008G", "2019MNRAS.485.4260S", "2019MNRAS.486.2494G", "2019PASP..131i4501K", "2019PhRvD..99a5013C", "2019PhRvD..99b4029D", "2019PhRvD..99b4048R", "2019PhRvD..99b4049C", "2019PhRvD..99h4022N", "2019PhRvD.100d3008S", "2019PhRvD.100d3011C", "2019PhRvD.100d3535D", "2019PhRvD.100f1101A", "2019PhRvD.100h3005D", "2019PhRvL.122l1101F", "2019PhRvX...9c1040A", "2019PrPNP.10903714B", "2019RPPh...82a6903C", "2019arXiv190500408A", "2020ApJ...900L..13A", "2020CQGra..37n5001D", "2020JCAP...05..050Y", "2020LRR....23....3A", "2020LRR....23....4B", "2020PhRvD.101b3016L", "2020PhRvD.101f4057G", "2020PhRvD.102l4037T", "2020arXiv200811316F", "2021A&A...649A.114G", "2021ApJ...914L..40D", "2021ApJ...921L..25L", "2021CQGra..38h5013U", "2021CQGra..38i5004A", "2021JPlPh..87a8402A", "2021MNRAS.500.4213M", "2021SoftX..1400677F", "2022A&A...659A..84A", "2022CQGra..39e5002A", "2022PASP..134f5001B", "2022PTEP.2022f3F01A", "2022PhRvD.105l3007F", "2022PhRvD.106l3529G", "2022PrPNP.12503948A", "2023A&A...674A..56R", "2023ApJ...952..157A", "2023IAUS..363...92C", "2023NatSR..1317706L", "2023PhRvD.108d2003H", "2024PrPNP.13404083C" ]
[ "astronomy" ]
14
[ "binaries: general", "gamma-ray burst: general", "gravitational waves", "stars: black holes", "stars: neutron", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
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[ "10.3847/2041-8205/832/2/L21", "10.48550/arXiv.1607.07456" ]
1607
1607.07456_arXiv.txt
16
7
1607.07456
We report here the non-detection of gravitational waves from the merger of binary-neutron star systems and neutron star-black hole systems during the first observing run of the Advanced Laser Interferometer Gravitational-wave Observatory (LIGO). In particular, we searched for gravitational-wave signals from binary-neutron star systems with component masses \in [1,3] {M}<SUB>⊙ </SUB> and component dimensionless spins &lt;0.05. We also searched for neutron star-black hole systems with the same neutron star parameters, black hole mass \in [2,99] {M}<SUB>⊙ </SUB>, and no restriction on the black hole spin magnitude. We assess the sensitivity of the two LIGO detectors to these systems and find that they could have detected the merger of binary-neutron star systems with component mass distributions of 1.35 ± 0.13 M <SUB>⊙</SUB> at a volume-weighted average distance of ∼70 Mpc, and for neutron star-black hole systems with neutron star masses of 1.4 M <SUB>⊙</SUB> and black hole masses of at least 5 M <SUB>⊙</SUB>, a volume-weighted average distance of at least ∼110 Mpc. From this we constrain with 90% confidence the merger rate to be less than 12,600 Gpc<SUP>-3</SUP> yr<SUP>-1</SUP> for binary-neutron star systems and less than 3600 Gpc<SUP>-3</SUP> yr<SUP>-1</SUP> for neutron star-black hole systems. We discuss the astrophysical implications of these results, which we find to be in conflict with only the most optimistic predictions. However, we find that if no detection of neutron star-binary mergers is made in the next two Advanced LIGO and Advanced Virgo observing runs we would place significant constraints on the merger rates. Finally, assuming a rate of {10}<SUB>-7</SUB><SUP>+20</SUP> Gpc<SUP>-3</SUP> yr<SUP>-1</SUP>, short gamma-ray bursts beamed toward the Earth, and assuming that all short gamma-ray bursts have binary-neutron star (neutron star-black hole) progenitors, we can use our 90% confidence rate upper limits to constrain the beaming angle of the gamma-ray burst to be greater than 2\buildrel{\circ}\over{.} {3}<SUB>-1.1</SUB><SUP>+1.7</SUP> (4\buildrel{\circ}\over{.} {3}<SUB>-1.9</SUB><SUP>+3.1</SUP>).
false
[ "neutron star masses", "neutron star", "black hole masses", "black hole", "neutron star-black hole systems", "binary-neutron star systems", "neutron star-binary mergers", "Advanced LIGO", "the same neutron star parameters", "Advanced Virgo", "(neutron star-black hole", "∼70 Mpc", "the black hole spin magnitude", "component masses", "component mass distributions", "LIGO", "mass \\in", "short gamma-ray bursts", "yr", "gravitational waves" ]
7.322193
2.449919
-1
12406239
[ "Rodenbeck, Kai", "Schleicher, Dominik R. G." ]
2016A&A...593A..89R
[ "Magnetic fields during galaxy mergers" ]
13
[ "Max-Planck-Institute for Solar System Research, Justus-von-Liebig-Weg 3, 37077, Göttingen, Germany", "Departamento de Astronomía, Facultad Ciencias Físicas y Matemáticas, Universidad de Concepción, Av. Esteban Iturra s/n Barrio Universitario, Casilla 160-C, Concepción, Chile" ]
[ "2016A&A...593A..77S", "2017PhDT.........3G", "2017PhRvE..95c3206G", "2018MNRAS.475.3283C", "2021MNRAS.506..229W", "2022arXiv220307241S", "2023ARA&A..61..561B", "2023ApJ...942L..13L", "2023EPJP..138..590S", "2023IAUS..362...94S", "2023arXiv231017036K", "2024arXiv240604242N", "2024arXiv240612532S" ]
[ "astronomy" ]
7
[ "Galaxy: evolution", "galaxies: interactions", "galaxies: magnetic fields", "methods: numerical", "Astrophysics - Astrophysics of Galaxies" ]
[ "1970ApJ...160..811F", "1977egsp.conf..401T", "1979JCoPh..32..101V", "1980Natur.285..643M", "1981A&A....96..111H", "1983ApJS...53..459C", "1985A&A...147L...6D", "1985AJ.....90..708K", "1985ApJ...298L...7H", "1985MNRAS.214...87J", "1987ApJ...320..238S", "1988ApJ...328...88S", "1988ApJ...331..699B", "1989MNRAS.240..329L", "1990A&A...236..333H", "1992ApJ...393..484B", "1992ApJ...400..460H", "1993A&A...269..581R", "1993AJ....105..864S", "1993AJ....105.1730C", "1993AJ....106.1337G", "1993ApJ...409..548H", "1995AJ....110..129K", "1995ApJ...453..100E", "1995ApJ...453..139E", "1996AJ....111..655H", "1996ARA&A..34..749S", "1996ApJ...464..641M", "1997A&A...320...54N", "1997A&A...325...81P", "1998ARA&A..36..189K", "1998ApJ...508L..43F", "1998MNRAS.297..143R", "1999ApJ...512L..99G", "1999Natur.397..324B", "2000A&A...355..128C", "2000MNRAS.318..124G", "2001ApJ...554..803Y", "2001ApJ...554.1035F", "2002AJ....123.1881C", "2002JCoPh.175..645D", "2002NJPh....4...84S", "2003AJ....126.2171G", "2004A&A...417..541C", "2004astro.ph..3044O", "2005A&A...444..739B", "2005PhR...417....1B", "2007AJ....134..118Z", "2007AJ....134.2124R", "2007MNRAS.377.1439B", "2008A&A...486L..35G", "2009ApJ...693.1449C", "2009ApJ...696...96W", "2009ApJ...700..358T", "2009ApJ...706..482M", "2009ApJ...706L.155H", "2009MNRAS.398.1082B", "2010A&A...518L..31I", "2010A&A...518L..44K", "2010AJ....139.1212S", "2010ApJ...712..536K", "2010ApJ...716.1438K", "2010ApJS..186..308C", "2010MNRAS.407..705W", "2010MNRAS.409...92J", "2010Natur.466.1082M", "2011A&A...529A..94C", "2011A&A...533A..22D", "2011ApJ...736..139K", "2011ApJ...739L..23H", "2011ApJS..192....9T", "2011MNRAS.410.1155B", "2011MNRAS.412.2396F", "2011MNRAS.414.2511V", "2011MNRAS.415.3189K", "2011PhRvL.107k4504F", "2012ApJ...745...65U", "2012ApJ...756..141B", "2012ApJ...758L..39T", "2012ApJ...761..140C", "2012ApJ...761..156F", "2012MNRAS.419.3571G", "2012PhRvE..85b6303S", "2013A&A...556A.142S", "2013A&A...557A.129T", "2013ApJ...771..120P", "2013MNRAS.429..967G", "2013MNRAS.432..176P", "2013MNRAS.434.2600F", "2013NJPh...15b3017S", "2013PhRvE..87e3110P", "2014A&A...566A..40M", "2014AJ....147..103H", "2014ApJ...783L..20P", "2014ApJ...791L..34B", "2014ApJS..211...19B", "2014MNRAS.440.1551L", "2014MNRAS.442.3407B", "2015A&A...578A..94M", "2015A&A...579A..32L", "2015MNRAS.447.2123C", "2015MNRAS.454.1545B", "2015NJPh...17b3070G", "2015PhRvE..92b3010S", "2016A&A...593A..77S", "2016ApJ...827..109S", "2016PhPl...23f2316V", "2016PhPl...23f2317G" ]
[ "10.1051/0004-6361/201527393", "10.48550/arXiv.1607.00871" ]
1607
1607.00871_arXiv.txt
Galaxy mergers are central for the evolution of galaxies, and have a strong influence on many of their properties. In a pioneering study to investigate the impact of mergers, \citet{Toomre77} has proposed a merging sequence of 11 peculiar galaxies spanning a range of pre- and post-mergers that became known as the {\em Toomre sequence}. This sequence suggested that the final product of mergers is likely to resemble an elliptical galaxy. This behavior was confirmed by numerical simulations following the N-body dynamics of the stellar population, finding the characteristic formation of $r^{1/4}$ profiles observed in elliptical galaxies \citep{Barnes88, Barnes92, Hernquist92, Hernquist93, Privon13}. Observational studies have provided evidence for this scenario, reporting a characteristic $r^{1/4}$ profile in merger remnants \citep{Joseph85} and finding low surface brightness loops and shells in elliptical galaxies, which are likely the remnants from previous mergers \citep{Malin80, Schweizer88}. \citet{Forbes98} have shown that late-stage disk-disk mergers and ellipticals with young stellar populations deviate from the fundamental plane of elliptical galaxies due to a centrally located starburst induced by the merger event. Gravitational instabilities acting during the merger event may channel significant gas flows towards the center of the galaxy, enhancing the activity of star formation or of a central supermassive black hole \citep{Mihos96}. High molecular gas fractions have been observed in IRAS starburst galaxies, which are considered as gas-rich systems close to the final stages of merging \citep{Sanders96, Planesas97, Kennicutt98}. The enhanced star-formation activity has been confirmed both in the far-infrared \citep{Kennicutt98} as well as in the radio \citep{Hummel81, Hummel90} and the optical \citep{Keel85}. Near-infrared observations with the Hubble Space Telescope show a distinct trend for the nuclei to become more luminous in more advanced merging stages \citep{Rossa07}. Similarly, the fraction of interacting systems in IRAS-selected samples increases with far-infrared luminosity \citep{Lawrence89, Gallimore93}. \citet{Gao99} have used the projected separation between the nuclei of merging galaxies as an estimator for the interaction stage, finding clear evidence for increasing star formation efficiencies, defined by the ratio of far-infrared luminosity to molecular hydrogen mass, with decreasing spatial separation. The latter potentially indicates more efficient star formation in higher-density gas. The increase of star formation activity during the merger event is also reflected in X-ray observations \citep{Read98, Brassington07} and in the UV \citep{Smith10}. \citet{Hibbard96} investigated the cold gas properties from a small sample of pre- and post-mergers, finding strong differences in the HI distribution in pre- and post-mergers, with increasing fractions of HI outside the optical bodies towards the later stages. They propose that the latter is due to the ejection of cold gas in tidal features. \citet{Georgakakis00} investigated a sample of $53$ interacting galaxy pairs, combining radio observations with optical data by \citet{Keel95}, far-infrared data by \citet{Gao99} and $60$~$\mu$m data from the sample of \citet{Surace93} based on the IRAS Bright Galaxy Sample \citep{Soifer87}. They defined the interaction parameter in terms of the effective age before or after the merger, confirming a significant increase of the star formation efficiency during the merger event. The latter is reflected in a correlation between the star formation efficiency and the H$_2$ surface density. The amount of molecular gas is found to increase during the merger as well, and to fall off after the merger. Beyond the statistical samples, some interacting galaxy pairs have been studied in particular detail. For instance, the Antennae galaxies (NGC~4038/39) have been explored with Chandra in the X-rays \citep{Fabbiano01}, with Herschel-PACS revealing obscured star formation \citep{Klaas10} and with the Submillimeter Array pursuing high-resolution CO observations \citep{Ueda12}. The change of the gas distribution due to the merger event may also be relevant for the cosmological growth of black holes \citep{Mayer10, Treister12, Ferrara13, Capelo15}. In addition to star formation and the molecular gas, magnetic fields are a relevant component of galaxies, which can influence the star formation rate \citep[e.g.][]{Banerjee09, Vazquez11, Federrath12}. The structure of the galactic magnetic field has been studied in detail in nearby galaxies \citep[e.g.][]{Beck99, Beck05, Fletcher11, Tabatabaei13, Heesen14}. More recently, the presence of magnetic fields was confirmed also in nearby dwarf galaxies \citep{Chyzy00, Kepley10, Heesen11, Kepley11, Chyzy11}. The far-infrared - radio correlation suggests that the presence of magnetic fields is strongly correlated with star formation activity, where the magnetic field strength $B$ scales approximately as the star formation surface density $\Sigma_{\rm SFR}$ to the one third \citep{deJong85, Helou85, Niklas97, Yun01, Basu12, Schleicher13c}, and has also been confirmed at high redshift \citep{Murphy09, Jarvis10, Ivison10, Bourne11, Casey12}. While radio data have been available for interacting galaxy pairs early on \citep[e.g.][]{Condon83, Condon93, Condon02}, the first detailed analysis of the magnetic field structure in an interacting galaxy pair has been pursued by \citet{Chyzy04} for the Antennae galaxies. With $\sim20$~$\mu$G, the mean total magnetic field strength is significantly greater than in isolated spirals, and could be enhanced via interaction-induced star formation. The field regularity appears to be lower compared to isolated galaxies, reflecting additional distortions or feedback related to star formation. The observed structures are roughly consistent with magneto-hydrodynamical (MHD) simulations by \citet{Kotarba10} based on a smoothed particle hydrodynamics (SPH) aproach. The galaxy pair NGC~2207/IC~2163 has been observed in the optical by \citet{Elmegreen95a}, and was modeled with N-body simulations by \citet{Elmegreen95b} reproducing the main dynamical features. The radio emission at $4.86$~GHz covers the optical extent of the system, but includes a hole in the central part of NGC~2207 \citep{Condon83}. For the Taffy (UGC~12914/15), a galaxy pair with a nearly head-on encounter, a significant amount of data is available, including radio and HI \citep{Condon93} as well as CO and the dust continuum \citep{Gao03, Zhu07}. A similar pair of post-collision spiral galaxies, UGC~813/16 was detected at $1.40$~GHz and forms the so-called Taffy2 \citep{Condon02}. The first more systematic study of magnetic fields in a complete merging sequence has been pursued by \citet{Drzazga11}. For this purpose, they compiled a sample of data from the Very Large Array (VLA) based on the merging sequence of \citet{Toomre77} and \citet{Brassington07} as well as other interacting system, leading to a total of $24$ interacting galaxy pairs, including pre- and post-merger stages. While they do not consider them as a complete sample of colliding galaxies, they do represent the different stages of tidal interaction and merging. For this sample, a correlation of the average magnetic field strength in the whole galaxy and in the central region is found with the star formation surface density $\Sigma_{\rm SFR}$. The central magnetic field strength scales as $\Sigma_{\rm SFR}^{0.27\pm0.03}$, consistent with the scaling relations in spiral galaxies \citep{Niklas97}, and in fact \citet{Drzazga11} report that their sample lies on the far-infrared - radio correlation for spiral galaxies. The evolution of magnetic fields during merger events is also interesting from a theoretical point of view, as the merger process leads to the production of turbulence, which can enhance both the small-scale magnetic fields via the small-scale dynamo \citep[e.g.][]{Scheko02, Cho09, FederrathPRL, Schober12a, Schleicher13b,Bhat14, Schober15}, and also affect the formation of a large-scale field via dynamo processes \citep{Ruediger93, Brandenburg05, Gressel13, Park13}. The number of numerical simulations exploring such merger scenarios is however still limited. Beyond the MHD-SPH simulations of the Antennae galaxies pursued by \citet{Kotarba10}, the impact of multiple mergers has been explored with the same approach by \citet{Kotarba11}. Magnetic field amplification during minor galaxy mergers has been investigated with a similar approach by \citet{Geng12}. Grid-based simulations of merging systems have been pursued in a 2D-approximation by \citet{Moss14}. However, as turbulence is essentially a 3D phenomenon, it is desireable to explore the evolution of magnetic fields during merger events with three-dimensional grid-based techniques. In this paper, we present a toy model to explore 3D effects during the merger of magnetized gaseous disks. {While this model still neglects the impact of the dark matter on the merger events, it shows that geometric projection events can be rather important when interpreting large-scale averages in observations of magnetic fields. We indeed show via a parameter study that the occurence of characteristic peaks due to projection effects is rather insensitive to the details of the orbit.} Our simulation setup is presented in section~\ref{setup} and the main results are given in section~\ref{results}. We show that our simulations can already reproduce some of the main features in the merger sequence by \citet{Drzazga11}. Our results and their implications for the interpretation of observational data are discussed in section~\ref{discussion}.
\label{discussion} To explore the evolution of magnetic fields during merger events, we present a toy model to simulate the evolution of magnetic fields during the mergers. For this purpose, we have modeled the interaction of two gaseous magnetized disks to explore how the interaction of the gas changes the structure of the magnetic field. In these idealized simulations, we have neglected the presence of a stellar component and a dark matter halo. {While the neglect of the dark matter halo can certainly have an impact on the orbit and change the dynamics, we have also seen in our sample of simulations that a set of characteristic features is rather insensitive to the details of the orbit. We describe in more detail below how part of the results may change after the inclusion of dark matter halos.} In spite of the simplicity of our model, we are able to reproduce central characteristics of the magnetic field evolution in the merger sequence provided by \citet{Drzazga11}. In particular, they have shown a characteristic peak in the average field strength of the galaxy, its central regions and also in the outer parts of the galaxy when plotted against the interaction parameter. For comparison, we have followed the time evolution of the magnetic field in the central $5$~kpc of the galaxy, within a $25$~kpc radius and in the region between $5$~kpc and $25$~kpc. We expect the latter to approximately resemble the three cases distinguished by \citet{Drzazga11}. {In our reference run}, we find two characteristic peaks in these quantities during the merger event. The peak in the average field strength within the galaxy and its outer region can be naturally explained by geometrical effects when the core of the second disk falls into the region covered by the $25$~kpc radius. Within the central $5$~kpc, both the magnetic field strength and the central density are actually enhanced, due to the increased angular momentum transport induced by the tidal torques within the configuration. The first peak occurs when the two cores for the first time come very close, but without merging, while the second peak occurs when the two cores are in fact about to merge. {We however note that the occurence of two peaks is not critical and indeed one peak in the averaged field strength, as observed in other simulations, is indeed enough to explain the sequence derived by \citet{Drzazga11}.} Despite their simplifications, our simulations can thus naturally explain the occurence of a peak in the magnetic field strength during the merger. We note that the time sequence produced in the simulation may not correspond one to one to the merger sequence produced by \citet{Drzazga11}, and in particular the interaction parameter may not uniquely correspond to a single time in the simulation, but rather to the characteristic times when the central cores come very close. Concerning the physical interpretation, our results show that one must carefully distinguish between geometrical effects on the one hand, as the average field strength will increase when the core of one galaxy enters the outer parts of another one{, and physical changes in the magnetic field strength in particular in the center of the galaxy, either due to density enhancements or potentially due to dynamo effects.} High-resolution observations to explore the central regions of merging galaxies are therefore particularly valuable to further explore the increase of the magnetic field strength during the merger along with its physical origin. \begin{figure*}[htbp] \begin{center} \includegraphics[scale=0.55]{C_545_rot_90_r_Bmag.eps} \includegraphics[scale=0.38]{C_545_rot_90_r_proj_MagneticPressure_z.eps} \caption{Comparison run with a relative disk orientation of $90^\circ$ {(run B)}. Top: Time evolution of the mean magnetic field strength on different scales. Bottom: Projection of the magnetic pressure at different times.} \label{perp} \end{center} \end{figure*} We have checked that the results presented here do not strongly depend on the parameters in the simulations, such as the mass of the galaxy, their relative velocity or the initial magnetic field strength. In some cases, we found the emergence of only one peak rather than two, depending on the detailed evolution of the final merger. {Such behavior tends to occur when the relative velocities are lower between the galaxies, leading to a more rapid merger event. In this case, there will be only one event where both disks fall into the observed beam, leading to two contributions. This also confirms that geometrical effects play the dominant role in case of large beams, and only high-resolution observations can help to assess the physical behavior of the magnetic fields.} {For realistic simulations, a dark matter halo should be included for both galaxies. To zero order, the latter will enhance both the gravitational and the inertial mass. While the baryons in the disks were strongly interacting leading to a relevant amount of friction, the latter will be reduced for the mass in dark matter. However, we note that even the dark matter halos will experience dynamical friction while moving through the circumgalactic material, including both baryonic and dark matter. While a precise orbit can only be obtained via dedicated simulations for each specific configuration, we expect that overall friction effects will be less relevant in a dark matter dominated regime, suggesting that the final merger may be somewhat delayed. The latter may affect the overall probability to find a specific galaxy in a certain merger stage, while the resulting geometrical projection effects are likely the same.} \begin{figure*}[htbp] \begin{center} \includegraphics[scale=0.55]{C_545_MR_03_Bmag.eps} \includegraphics[scale=0.38]{C_545_MR_03_proj_MagneticPressure_z.eps} \caption{Comparison run with a mass ratio of $1:3$ {(run C)}. Top: Time evolution of the mean magnetic field strength on different scales. Bottom: Projection of the magnetic pressure at different times.} \label{ratio} \end{center} \end{figure*} As we noted in the introduction, previous investigations concerning the evolution of magnetic fields during galaxy mergers have been predominantly pursued using MHD-SPH simulations \citep[e.g.][]{Kotarba10, Kotarba11, Geng12}, or in the 2D-approximation using a grid-based approach \citep{Moss14}. Even the number of studies of isolated galaxies including magnetic fields is still limited. In this respect, we note the simulations pursued by \citet{Wang09} exploring galaxy formation in a cosmological context, where they neglected stellar feedback. More recently, \citet{Pakmor13} pursued simulations of isolated disk galaxies including magnetic fields and stellar feedback using a moving-mesh approach. Magnetic field amplification in isolated galaxies is further explored in the 2D approximation \citep{Moss15} or through the simulation of local volumes \citep{Gressel08}. The interaction of the dynamo process with cosmic rays has further been explored by \citet{Hanasz09}. Hydrodynamical simulations of isolated galaxies have been pursued with the Enzo code by \citet[e.g.][]{Tasker09}, and gravitational instabilities in isolated disks have been explored by \citet{Lichtenberg15}. While we have shown here that even simplified models can already reproduce some of the main features in the observed merger sequence, future simulations should aim for a more detailed and more realistic modeling, accounting for the effects of feedback, and employing cosmological initial conditions. In the context of isolated disk galaxies, the latter has been achieved by \citet{Wang09}, \citet{Latif14} and \citet{Pakmor14}, but it is still outstanding in the context of galaxy mergers. An additional problem is the limited resolution in numerical simulations, implying that galaxy-scale simulations will never be able to resolve turbulence and small-scale magnetic fields on all relevant scales. In the magneto-hydrodynamical equations, it may therefore be necessary to adopt nonlinear closures to explore the resulting dynamical implications \citep{Grete15}. The latter was shown to be relevant in hydrodynamical simulations of isolated galaxies \citep[e.g.][]{Braun14, Braun15}, and we expect similar effects to occur in MHD simulations.
16
7
1607.00871
Galaxy mergers are expected to play a central role for the evolution of galaxies and may have a strong effect on their magnetic fields. We present the first grid-based 3D magnetohydrodynamical simulations investigating the evolution of magnetic fields during merger events. For this purpose, we employed a simplified model considering the merger event of magnetized gaseous disks in the absence of stellar feedback and without a stellar or dark matter component. We show that our model naturally leads to the production of two peaks in the evolution of the average magnetic field strength within 5 kpc, within 25 kpc, and on scales in between 5 and 25 kpc. The latter is consistent with the peak in the magnetic field strength previously reported in a merger sequence of observed galaxies. We show that the peak on the galactic scale and in the outer regions is most likely due to geometrical effects, as the core of one galaxy enters the outskirts of the other one. In addition, the magnetic field within the central ~5 kpc is physically enhanced, which reflects the enhancement in density that is due to efficient angular momentum transport. We conclude that high-resolution observations of the central regions will be particularly relevant for probing the evolution of magnetic field structures during merger events.
false
[ "magnetic fields", "magnetic field structures", "merger events", "efficient angular momentum transport", "observed galaxies", "stellar feedback", "galaxies", "magnetized gaseous disks", "geometrical effects", "Galaxy mergers", "the average magnetic field strength", "the magnetic field strength", "the magnetic field", "their magnetic fields", "the central ~5 kpc", "scales", "a stellar or dark matter component", "density", "the merger event", "a merger sequence" ]
12.893283
9.438962
-1
12464338
[ "McLaughlin, J. A." ]
2016JApA...37....2M
[ "Behaviour of Magnetoacoustic Waves in the Neighbourhood of a Two-Dimensional Null Point: Initially Cylindrically Symmetric Perturbations" ]
2
[ "Department of Mathematics and Information Sciences, Northumbria University, Newcastle-upon-Tyne, United Kingdom" ]
[ "2019A&A...624A..90P", "2019A&A...629A..60H" ]
[ "astronomy" ]
3
[ "Magnetohydrodynamics (MHD)", "waves", "magnetic fields", "Sun: atmosphere", "corona", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1932hydr.book.....L", "1958ApJ...127..459F", "1978amms.book.....B", "1982A&A...112...16Z", "1983A&A...117..220H", "1991ApJ...376..355P", "1997ApJ...486L..67C", "1997SoPh..175...93N", "1998A&A...337..287E", "1998A&A...339..208B", "2001A&A...367..339B", "2001ApJ...548..473C", "2002A&A...392..319H", "2002ApJ...564..508R", "2002RSPSA.458.2307H", "2002SoPh..209..333B", "2003A&A...400.1065O", "2003ApJ...599..626B", "2004A&A...420.1129M", "2004A&A...425..741D", "2004SoPh..225...21C", "2004prma.book.....G", "2005A&A...435..313M", "2005A&A...436L..35O", "2006A&A...452..603M", "2006A&A...459..641M", "2006JPlPh..72..571M", "2006SoPh..237..143D", "2007JGRA..112.3103P", "2007PhPl...14e2106P", "2007Sci...317.1192T", "2008A&A...484L..47R", "2008SoPh..251..563M", "2009A&A...493..227M", "2009JPlPh..75..203A", "2009SoPh..254...51L", "2011A&A...527A.149M", "2011A&A...529A..20G", "2011A&A...533A..18M", "2011A&A...534A...2G", "2011PhPl...18h4506K", "2011SSRv..158..205M", "2012A&A...545A...9T", "2013A&A...558A.127T", "2013JApA...34..223M", "2013MNRAS.428...40C", "2013SoPh..288..205T", "2015A&A...577A.138K" ]
[ "10.1007/s12036-016-9376-y", "10.48550/arXiv.1607.02363" ]
1607
1607.02363_arXiv.txt
MHD waves are ubiquitous in the Sun's atmosphere (e.g. Tomczyk {\emph{et al.}} \citeyear{Tomczyk2007}) and a variety of observations have now demonstrated the existence of wave activity for the three fundamental MHD wave modes: namely Alfv\'en waves and fast and slow magnetoacoustic waves. Non-thermal line broadening and narrowing due to Alfv\'en waves has been reported by various authors, including Banerjee {\emph{et al.}} (\citeyear{Banerjee1998}), Erd\'elyi {\emph{et al.}} (\citeyear{Erdelyi1998}), Harrison {\emph{et al.}} (\citeyear{Harrison2002}) and O'Shea {\emph{et al.}} (\citeyear{OShea2003}; \citeyear{OShea2005}) and investigated both analytically (e.g. Dwivedi \& Srivastava \citeyear{DS2006}) and numerically (e.g. Chmielewski {\emph{et al.}} {\citeyear{Chmielewski2013}}, and references therein). MHD wave behaviour is influenced strongly by the underlying magnetic structure (topology) and so it is useful to look at the topology itself. Potential field extrapolations of the coronal magnetic field can be made from photospheric magnetograms and such extrapolations show the existence of two key features of the magnetic topology: {\emph{magentic null points}} and {\emph{separatrices}}. {{Null points}} are weaknesses in the magnetic field at which the field strength, and thus the Alfv\'en speed, is zero. {{Separatrices}} are topological features that separate regions of different magnetic connectivity and are an inevitable consequence of the isolated magnetic flux fragments in the photosphere. Detailed investigations of the coronal magnetic field, using such potential field calculations, can be found in Beveridge {\it{et al.}} (\citeyear{Beveridge2002}) and Brown \& Priest (\citeyear{Brown2001}). The number of resultant null points depends upon the complexity of the magnetic flux distribution and tens of thousands are estimated to be present (see, e.g., Close {\emph{et al.}} \citeyear{Close2004}; Longcope \citeyear{L2005}; R{\'e}gnier {\emph{et al.}} \citeyear{RPH2008}; Longcope \& Parnell \citeyear{LP2009}). MHD waves and magnetic topology {\emph{will}} encounter each other in the solar corona, e.g. waves emanating from a flare or CME will at some point encounter a coronal null point. MHD wave propagation within an inhomogeneous magnetic medium is a fundamental plasma process and the study of MHD wave behaviour in the neighbourhood of magnetic null points directly contributes to this area; see McLaughlin {\emph{et al.}} (\citeyear{McLaughlinREVIEW}) for a comprehensive review of the topic. The behaviour of linear MHD waves, both magnetoacoustic waves and Alfv\'en waves, has been investigated in the neighbourhood of a variety of 2D null points (e.g. McLaughlin \& Hood \citeyear{MH2004}; \citeyear{MH2005}; \citeyear{MH2006a}; \citeyear{MH2006b}; McLaughlin \citeyear{M2013}). Nonlinear and 3D MHD wave activity about coronal null points has also been investigated (e.g. Galsgaard {\emph{et al.}} \citeyear{Galsgaard2003}; Pontin \& {Galsgaard} \citeyear{PG2007}; Pontin {\emph{et al.}} \citeyear{PBG2007}; McLaughlin {\emph{et al.}} \citeyear{MFH2008}; \citeyear{McLaughlin2009}; Galsgaard \& Pontin \citeyear{klaus2011a}; \citeyear{klaus2011b}; Thurgood \& McLaughlin \citeyear{Thurgood2012}; \citeyear{Thurgood2013a}; \citeyear{Thurgood2013b}). {{ Authors have also considered an X-point magnetic field configuration with a longitudinal (along the X-line) magnetic field $B_\parallel$. This has the effect that now the fast magnetoacoustic wave and Alfv\'en wave are linearly coupled by the gradients in the field. McClements {\it{et al.}} (\citeyear{McClements2006}) investigated such a coupling with a weak longitudinal guide field present ($B_\parallel \ll B_\perp$) and Ben Ayed {\it{et al.}} (\citeyear{BenAyed2009}) considered a strong guide-field ($B_\parallel \gg B_\perp$). These authors found that the Alfv\'en wave is coupled into the fast mode, with the coupling strongest on the separatrices and far from the X-line. In the limit of $B_\parallel \to 0$, the two modes are decoupled and the results of 2D work are recovered. More recently, {Ku{\'z}ma} {\it{et al.}} (\citeyear{KMS2015}) investigated similar coupling for a X-line formed above two magnetic arcades, but now embedded in a model solar atmosphere with a realistic temperature distribution. They found that the formation of the Alfv\'en waves at the initial phase of temporal evolution is followed by linear coupling between Alfv\'en and magnetoacoustic waves at a later time. The Alfv\'en waves also experience phase mixing and scattering from inhomogeneous regions of Alfv\'en speed, and partial reflection from the model transition region. }} It is also clear that the plasma-$\beta$, i.e. the ratio of thermal plasma pressure to magnetic pressure, plays a key role. A very detailed and comprehensive set of 2D numerical simulations of wave propagation in a stratified magneto-atmosphere was conducted by Rosenthal {\it{et al.}} (\citeyear{Rosenthal2002}) and Bogdan {\it{et al.}} (\citeyear{Bogdan2003}). In these simulations, an oscillating piston generated both fast and slow MHD waves on a lower boundary and sent these waves up into the stratified magnetized plasma. Their calculations showed there was coupling between the fast and slow waves, and that this coupling was confined to a thin layer where the sound speed and the Alfv\'en velocity are comparable in magnitude, i.e. where the plasma-$\beta$ approaches unity. Away from this conversion zone, the waves were decoupled as either the magnetic pressure or thermal plasma pressure dominated. One of the aims of these papers was to see how the topology affected the propagation of waves, with the ratio of the sound speed to the Alfv\'en speed varying along every magnetic line of force. In this, their work and ours have the same ultimately goal; a fully 3D understanding of MHD wave propagation in the solar corona. Other authors have also looked at MHD mode coupling: Cally \& Bogdan (\citeyear{Cally1997}) describes 2D simulations in which both $f$-modes and $p$-modes are partially converted to slow magnetoacoustic gravity waves due to strong gravitational stratification. De Moortel {\it{et al.}} (2004) investigated driving slow waves on the boundary of a 2D geometry with a horizontal density variation, where they found coupling between slow and fast waves and phase mixing of the slow waves. The coupling of different wave modes has also been investigated by {Ferraro} \& {Plumpton} (\citeyear{Ferraro1958}), with Meijer G-functions by {Zhugzhd} \& {Dzhalilov} (\citeyear{Zhugzhd}), and with hypergeometric $_2F_3$ functions by Cally (\citeyear{Cally2001}). All these works considered mode coupling through a gravitational stratification, i.e. a vertical density inhomogenity. Finally, the coupling of fast waves and Alfv\'en waves has been investigated by Parker (\citeyear{Parker1991}) for linear MHD with a density gradient and by Nakariakov {\it{et al.}} (\citeyear{Valery1997}) for nonlinear excitation. In this paper, we will investigate the behaviour of magnetoacoustic waves within inhomogeneous magnetic media. We will concentrate our investigations on wave behaviour excited via initially cylindrical-symmetric perturbations. Our paper has two aims: Firstly, we will investigate the behaviour of (fast) magnetoacoustic waves in a $\beta=0$ plasma using numerical, analytical and semi-analytical techniques. Secondly, we lift the $\beta=0$ assumption and study a $\beta \neq 0$ plasma. This naturally introduces slow waves to the system and so we will investigate the behaviour of both types of magnetoacoustic waves around a null point. Two papers are key to our investigation: Firstly, McLaughlin \& Hood (\citeyear{MH2004}) investigated the behaviour of the fast magnetoacoustic wave in a $\beta=0$ plasma within a Cartesian geometry. They found that the fast magnetoacoustic wave was attracted to the null via a refraction effect and that all the wave energy accumulated at the null. Secondly, McLaughlin \& Hood (\citeyear{MH2006b}) extended the 2004 model to include plasma pressure in a $\beta \neq 0$ system. This led to the introduction of slow magnetoacoustic waves and coupling between the two types of magnetoacoustic waves. However, the resultant behaviour was extremely complex and the investigate was again limited to a Cartesian geometry. In this paper, we will investigate the behaviour of magnetoacoustic waves in a $\beta \neq 0$ plasma excited via initially cylindrical-symmetric perturbations. It is hoped that our results will help begin to explain the complex resultant behaviour observed in McLaughlin \& Hood (\citeyear{MH2006b}) and hence contribute to the overall understanding of MHD mode conversion across the $\beta=1$ layer. Our paper has the following outline: The basic setup, equations and assumptions are described in $\S\ref{sec:1.1}$. The analytical, numerical and semi-analytical results for a $\beta=0$ plasma are presented in $\S\ref{sec:2.3.3}$ and the numerical results for a $\beta \neq 0$ plasma appear in $\S\ref{polarcoorindatesbetaneq0}$. The discussions and conclusions are given in $\S\ref{sec:6.10}$.
In $\S\ref{sec:2.3.3}$, we investigated the behaviour of an initially cylindrically-symmetric fast magnetoacoustic wave around a 2D null point under the $\beta=0$ assumption. Using polar coordinates, we derived a governing wave equation with a spatially-varying characteristic speed (the Alfv\'en speed) and we solved this equation analytically by deriving a {\emph{Klein-Gordon}} equation and then solving separately for $m=0$, which led to a D'Alembert-type solution, and $m\neq 0$ which led to a Bessel-type solution (equation \ref{solutuion}). It is interesting to note that solution (\ref{solutuion}) is only valid for $s \geq 0$, i.e. $ t \geq \pm \ln {\frac {r} {r_0} }$, and that the same final result is gained from substituting $s = \sqrt {u^2 - t^2 }$ or $s = \sqrt {t^2 - u^2 }$, since $J_0(s) = J_0(-s) $. We can interpret this as follows: if we consider the boundary of our system to be a shell at radius $r_0$, we can interpret the $\pm$ ambiguity on $u$ as a boundary disturbance splitting into two waves; one propagating outwards ($r$ increasing so $r>r_0$, i.e. $u = \ln{\frac{r}{r_0}}$ solution) and one propagating inwards ($r$ decreasing so $r<r_0$, i.e. $u = -\ln{\frac{r}{r_0}}$). Note that the inequality on $r$ here dictates the flow of information; the perturbation starts on the boundary and there is no disturbance in front of the wave, i.e. the inequality that restricts $r$ from taking certain values until time has elapsed is interpreted as regions in the system not yet affected by the perturbation; as the information has not yet had the time to reach there since the wavefront propagates at a finite speed. Thus, if we are interested in the region inside $r=r_0$ including the origin (which is the location of our null) then we are interested primarily in the substitution $u = - \ln { \frac {r} {r_0} } $, with $r$ starting at $r_0$ and decreasing as $t$ evolves. We also solved the $\beta=0$ governing wave equation using numerical techniques in $\S\ref{sec:2.3.3.a}$. We find that the linear, $\beta=0$ fast magnetoacoustic wave splits into two waves; one approaching the null and the other propagating away from it. The wave propagating towards the null has the shape of an annulus. We find that this annulus contracts, but keeps its original ratios (distance between the leading-and-middle wavefronts compared to middle-and-trailing wavefronts). This was seen in Figure \ref{fig:2.3.3.2}. Since the Alfv\'en speed is spatially varying (i.e. $\sim r$, see equation \ref{fastbetapolar}), a \emph{refraction} effect focuses the wave into the null point. This is the same refraction effect found in McLaughlin \& Hood (\citeyear{MH2004}). Finally, we investigated our system using a semi-analytical WKB approach in $\S\ref{sec:2.3.3.c}$. This can be seen in Figure \ref{fig:2.3.3.3}. As expected, the agreement between Figures \ref{fig:2.3.3.2} and \ref{fig:2.3.3.3} is excellent; the semi-analytical WKB and numerical solutions lie on top of each other. We can also see in Figure \ref{fig:2.3.3.3} how the ratio between the leading-and-middle and between the middle-and-trailing of the pulse is preserved. The wave focuses on the null point and contracts around it. In addition, equations (\ref{polar_zero_characteristics}) can be solved analytically by forming \begin{eqnarray*} { \frac{dp}{ds} } \;\big{/} \;{ \frac{dr}{ds} }= \frac{dp}{dr} = -\frac{p}{r} \quad \Longrightarrow \quad \log{r} = -\log{p} + {\rm{constant}} \quad \Longrightarrow \quad rp = -\omega \nonumber \end{eqnarray*} and so \begin{eqnarray} \frac {dp}{ds} = \omega p \; , \quad \frac {dr}{ds} = -\omega r\; , \quad p = -\frac {\omega}{r_0} e^{\omega s} \; , \quad r = {r_0} e^{-\omega s} \label{evolution}\;, \end{eqnarray} where the initial conditions dictate the constants of integration. From equations (\ref{evolution}) we see $r = {r_0} e^{-\omega s}$ and so the wave, which focuses on the null and contracts around it, never actually reaches the null in a finite time, due to the exponential decay of $r$. \subsection{$\beta \neq 0$ plasma} In $\S\ref{polarcoorindatesbetaneq0}$, we investigated the behaviour of an initially cylindrically-symmetric fast magnetoacoustic wave around a 2D null point in a $\beta \neq 0$ plasma. This can be seen in Figure \ref{fig:2.3.3.2_MAGIC}. We find that the fast wave split into two waves; one approaching the origin and the other travelling away from it; as expected. The wave propagating towards the origin initially has the shape of an annulus. We find that the annulus contracts (as in $\S\ref{sec:2.3.3.a}$ and Figure \ref{fig:2.3.3.2_MAGIC}) and that, at least initially, this contraction appeared to preserve the original ratios (distance between the leading-and-middle wavefronts compared to middle-and-trailing wavefronts). However, as the wave continues to propagate towards the origin, it is distorted significantly from its original shape: there is a decrease in overall wave speed along the $x=0$ and $z=0$ axes (the separatrices) and so the annulus starts to take on a quasi-diamond shape; with the corners located along the separatrices. This can be seen in the second and third row of subfigures of Figure \ref{fig:2.3.3.2_MAGIC}. Eventually, the wave crosses the $c_s^2=v_A^2\:$ layer (indicated by a black circle in the figure, located at $r=0.456$ for $\beta_0=0.25$) where it begins a more complicated evolution: unlike that seen in the equivalent $\beta=0$ case in Figure \ref{fig:2.3.3.2_MAGIC}. The formation of the quasi-diamond shape in Figure \ref{fig:2.3.3.2_MAGIC} is due to a decrease in the overall wave speed along the separatrices. This decrease is wave speed can be understood by investigating the perturbed pressure, $p_1$, and this can be seen in Figure \ref{fig:2.3.3.4}. We see that $p_1$ propagates towards the null similar to the propagation of the fast wave and is zero along the axes, i.e. the lines $x=0$ and $z=0$. Hence, because of the alternating nature of the pressure, the maximum gradients in pressure will occur along these locations, i.e. {\emph{along the separatrices}}. This pressure gradient acts against the magnetic forces in the momentum equation and thus reduces the acceleration of the fast wave along the separatrices, i.e. the magnitude of $\frac {\partial}{\partial t} {\rm{v}_\perp}$ is smaller along the separatrices leading to the deceleration as seen in Figure \ref{fig:2.3.3.2_MAGIC}. Note also that the pressure is increasing all the time and this can be seen in Figure \ref{fig:2.3.3.5}. Note that in this paper we do not describe the evolution of ${\rm{v}_\perp}$ after it crosses the $c_s^2=v_A^2\:$ layer; this crossing occurs at approximately $t=1.5$. As the wave crosses the $c_s^2=v_A^2\:$ layer, complex MHD mode conversion occurs. However, the description of such mode conversion is not the focus of this current paper and the resultant mode conversion has already been reported by McLaughlin \& Hood (\citeyear{MH2006b}). Instead, this paper focuses on (i) the nature of the wave propagation {\emph{before}} crossing the $c_s^2=v_A^2\:$ layer and (ii) comparing and contrasting this behaviour to that seen in the $\beta=0$ system. Thus, our main conclusion for the $\beta \neq 0$ system is related to the explanation of the quasi-diamond shape in Figure \ref{fig:2.3.3.2_MAGIC} and that this deformation in wave morphology was absent in the $\beta=0$ set-up. Note that McLaughlin \& Hood (\citeyear{MH2006b}) {\emph{does not}} include our insights related to the formation of the quasi-diamond pattern as well as its explanation in terms of the maximum gradients in pressure occurring along the separatrices. We also note that early on in its evolution, the $\beta \neq 0$ fast wave evolves in a similar manner to its $\beta=0$ equivalent. By looking at the equations (\ref{finitebetaequations}), we see this makes sense; at large radii the pressure terms are negligible and so the Alfv\'en speed is essentially spatially varying like $r$, and so the refraction effect dominates the evolution. \begin{figure}[t] \begin{center} \includegraphics[width=6.0in]{Figure8_LOWER_RESOLUTION.eps} \caption{Contours of numerical simulation of $p_1$ for a fast wave pulse initially located about a radius $\sqrt{x^2+z^2}=3$, and its resultant propagation at times $(a)$ $t$=0.2, $(b)$ $t$=0.6, $(c)$ $t$=1.0, $(d)$ $t=$1.4, $(e)$ $t$=1.8, $(f)$ $t$=2.2, labelling from top left to bottom right. The black circle indicates the position of the $c_s^2=v_A^2\:$ layer and the cross denotes the null point in the magnetic configuration. $p_1$ has an alternating form, where orange represents $p_1>0$ and blue represents $p_1<0$. The pressure appears to follow a $\sin{2\theta}$ pattern.} \label{fig:2.3.3.4} \end{center} \end{figure} \begin{figure}[h] \begin{center} \includegraphics[width=2.0in]{Figure7.eps} \caption{Increase in pressure as fast wave approaches and crosses the $c_s^2=v_A^2\:$ layer. The wave crosses the $c_s^2=v_A^2\:$ layer at approximately $t=1.5$.} \label{fig:2.3.3.5} \end{center} \end{figure} We can also looked at the behaviour of ${\rm{v}_\parallel}$ in $\S\ref{ein2}$ and this can be seen in Figure \ref{fig:2.3.3.9}. We observe that the wave has an alternating structure in the $\theta$-direction, i.e. positive and negative parts to the wave, unlike ${\rm{v}_\perp}$ which was always positive, and we note that the ratio ${\rm{v}_\parallel}$ : ${\rm{v}_\perp}$ of the amplitude of disturbances is $\beta_0$ : $1$. We also note that we have set an initial condition in ${\rm{v}}_\perp$ only: the initial condition on the parallel wave was ${\rm{v}_\parallel} (x,z,0) = 0$ in equations (\ref{initalconditionsIC}). Hence, the ${\rm{v}_\parallel}$ wave we are observing has been generated as a consequence of our ${\rm{v}}_\perp$ initial condition. By looking at equations (\ref{finitebetaequations}) and our initial conditions, we see that ${\rm{v}}_\perp$ acts as a driver or forcing term for ${\rm{v}_\parallel}$. Thus, we are solving the equivalent of a second-order differential equation with a forcing term, which is an inhomogeneous equation. The general solution to such equations consists of two parts; a {\emph{complementary function}} and a {\emph{particular integral}}. The complementary function is a solution to the corresponding homogeneous differential equation whereas the particular integral is a solution to the inhomogeneous differential equation. Hence, returning our attention to Figure \ref{fig:2.3.3.9}, we see that there should be two parts to the ${\rm{v}_\parallel}$ wave. We do see a part which has the same speed and frequency as the perpendicular component wave and, using the definition above, this wave can be thought of as the particular integral to the equations. There is also be a complementary function part to the wave, though it is difficult to see in the figure. As a consequence of our results from $\S\ref{sec:2.3.3}$ and $\S\ref{polarcoorindatesbetaneq0}$, we shall now explain how we interprete the waves seen in the perpendicular and parallel velocities. \subsection{Interpretating the waves we see in ${\bf{v}}_\perp$ and ${\bf{v}}_\parallel$} Our MHD system contains three key velocities: ${\bf{v}}_{\rm{Alfv{\acute{e}}n}}$, ${\bf{v}}_{\rm{slow}}$ and ${\bf{v}}_{\rm{fast}}$, that are all orthogonal and thus we may consider them as an {\emph{orthogonal basis}} of vectors for our system. In this paper, we do not consider the Alfv\'en wave, ${\bf{v}}_{\rm{Alfv{\acute{e}}n}}=v_y {\hat{\bf{y}}}$, and so our 2D vectors may be described in terms of the vectors ${\bf{v}}_{\rm{fast}}$ and ${\bf{v}}_{\rm{slow}}$. Due to our choice of coordinate system ($\S\ref{sec:1.2b}$) we choose to work in the directions perpendicular and parallel to the magnetic field. Thus, we may represent these two vectors in terms of ${\bf{v}}_{\rm{fast}}$ and ${\bf{v}}_{\rm{slow}}$, namely \begin{eqnarray*} {\bf{v}}_\perp = {A} {\bf{v}}_{\rm{fast}} + {B}{\bf{v}}_{\rm{slow}}\;\; ,\quad {\bf{v}}_\parallel = {C} {\bf{v}}_{\rm{fast}} + {D}{\bf{v}}_{\rm{slow}} \;. \end{eqnarray*} Alternatively, we may express our two magnetoacoustic velocities in terms of ${\bf{v}}_\perp$ and ${\bf{v}}_\parallel$, namely \begin{eqnarray*} {\bf{v}}_{\rm{fast}} = E {\bf{v}}_\perp + F {\bf{v}}_\parallel \;\; ,\quad {\bf{v}}_{\rm{slow}} =G {\bf{v}}_\perp + H {\bf{v}}_\parallel \;, \end{eqnarray*} where $A$, $B$, $C$, $D$, $E$, $F$, $G$ and $H$ are unknown functions that depend upon the magnetic geometry and (possibly) the plasma-$\beta$. This representation is only possible because both ${\bf{v}}_{\rm{fast}}$ and ${\bf{v}}_{\rm{slow}}$ and ${\bf{v}}_\perp$ and ${\bf{v}}_\parallel$ form orthogonal bases. However, we must be cautious: the concepts of fast and slow waves were originally derived for a uni-directional magnetic field and so these ideas may not carry over to more complex magnetic geometries quite as simply as claimed here. However, we recommend still utilizing terminology such as fast and slow wave in the interpretation of the waves in complex topologies, as well as the intuition gained from the uni-directional magnetic field models. We believe that a good way of interpreting the waves we see in our magnetic configuration is as follows \begin{eqnarray*} \rm{fast \;wave }&=&\rm{\;(large\; perpendicular \;component) \;}+\rm{\; (parallel \;component) } \nonumber\\ \rm{ }&=&\rm{\;(large\; component\; in\; {{v}_\perp})\; }+\rm{\; (component \;in\; {{v}_\parallel}) }\; , \nonumber\\ \rm{slow \;wave\; }&=&\rm{\; (small\; perpendicular \;component) \;}+\rm{ \;(parallel\; component)} \nonumber\\ \rm{ }&=&\rm{\;(small\; component\; in\; {{v}_\perp})\; }+\rm{\;(component\; in \;{{v}_\parallel}) }\;.\label{interpretationslowfast} \end{eqnarray*} In addition, our system consists of a region of low-$\beta$ plasma outside the $\beta=1$ layer and a region of high-$\beta$ plasma within; see Figure \ref{fig:areasofhighandlow}. This is understood from our definition of the plasma-$\beta$ for this magnetic field; $\beta \propto \left( {x^2+z^2} \right) ^{-1}$. Recall that slow and fast waves have differing properties depending on whether they are in a high or low-$\beta$ environment. To summarise: \vspace{0.25in} \begin{table}[h]\label{FSprop} \begin{center} \begin{tabular}{|c|c|c|} \hline & Fast Wave & Slow Wave\\ \hline\hline High-$\beta$ & \begin{tabular}{c} Behaves like sound wave\\ (speed $c_s$) \end{tabular} & \begin{tabular}{c} Guided along ${\bf{B}}_0$ \\ Transverse wave travelling at $v_A$ \end{tabular}\\ \hline Low-$\beta$ & \begin{tabular}{c} Propagates roughly isotropically \\ (speed $v_A$) \end{tabular} & \begin{tabular}{c} Guided along ${\bf{B}}_0$\\ Longitudinal wave travelling at speed $c_T$\\ \end{tabular}\\ \hline \end{tabular} \end{center} \end{table} In our investigations, we have sent a wave pulse into our system from a particular radius, i.e. in the low-$\beta$ region. At some point this wave has crossed the $\beta=1$ layer and entered the high-$\beta$ environment. Thus, we have a low-$\beta$ wave approaching the layer, coupling and mixing inside the layer and emerging as a mixture of high-$\beta$ fast and slow waves. We are driving waves in the perpendicular velocity component in a low-$\beta$ region (see Figure \ref{fig:areasofhighandlow}) and so we interpret this as predominantly low-$\beta$ fast wave. At this time there does not exist a robust set of rules connecting low and high-$\beta$ waves across the $\beta=1$ layer. It is hoped that the work presented here will help contribute to such a set of rules, specifically in what happens when a low-$\beta$ fast wave crosses the $\beta=1$ layer and becomes part high-$\beta$ fast wave and part high-$\beta$ slow wave. {{ We conclude that in a $\beta=0$ plasma, the fast wave cannot cross the null point and all the wave energy accumulates at that location. Thus, {\emph{null points will be locations for preferential heating from fast magnetoacoustic waves}}. For $\beta \neq 0$, the evolution is more complex and the fast wave now couples with the slow wave close to the $\beta=1$ layer. The resultant behaviour is controlled by the parameter $\beta_0$. Finally, there is as yet no unambiguous observational evidence for MHD wave behaviour in the vicinity of coronal null points. The successful detection of MHD waves around coronal null points will require advancements in two areas: high-spatial and high-temporal resolution imaging data as well as magnetic extrapolations from co-temporal magnetograms. Future missions, such as the {\it{Daniel K. Inouye Solar Telescope}} and {\it{Solar Orbiter}} may satisfy these requirements, and so the first detection of MHD waves in the neighbourhood of null points may be reported in the near future. }} \bigskip {\bf Acknowledgements} {The author acknowledges IDL support provided by STFC. JM wishes to thank Alan Hood for insightful discussions and constant encouragement.}
16
7
1607.02363
The propagation of magnetoacoustic waves in the neighbourhood of a 2D null point is investigated for both β=0 and β ≠ 0 plasmas. Previous work has shown that the Alfvén speed, here v <SUB> A </SUB>∝ r, plays a vital role in such systems and so a natural choice is to switch to polar coordinates. For β=0 plasma, we derive an analytical solution for the behaviour of the fast magnetoacoustic wave in terms of the Klein-Gordon equation. We also solve the system with a semi-analytical WKB approximation which shows that the β=0 wave focuses on the null and contracts around it but, due to exponential decay, never reaches the null in a finite time. For the β ≠ 0 plasma, we solve the system numerically and find the behaviour to be similar to that of the β=0 system at large radii, but completely different close to the null. We show that for an initially cylindrically-symmetric fast magnetoacoustic wave perturbation, there is a decrease in wave speed along the separatrices and so the perturbation starts to take on a quasi-diamond shape; with the corners located along the separatrices. This is due to the growth in pressure gradients that reach a maximum along the separatrices, which in turn reduces the acceleration of the fast wave along the separatrices leading to a deformation of the wave morphology.
false
[ "wave speed", "magnetoacoustic waves", "β=0 plasma", "polar coordinates", "such systems", "exponential decay", "the fast magnetoacoustic wave", "the fast wave", "the wave morphology", "a 2D null point", "large radii", "a quasi-diamond shape", "a semi-analytical WKB approximation", "turn", "contracts", "the β=0 system", "the separatrices", "SUB", "pressure gradients", "WKB" ]
12.439307
14.62079
2
12409306
[ "van der Marel, N.", "Cazzoletti, P.", "Pinilla, P.", "Garufi, A." ]
2016ApJ...832..178V
[ "Vortices and Spirals in the HD135344B Transition Disk" ]
72
[ "Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, 96822 Honolulu, HI, USA", "Max-Planck-Institut fur Extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany", "Leiden Observatory, P.O. Box 9513, 2300 RA Leiden, the Netherlands", "ETH, Zurich, Wolfgang-Pauli-Strasse 27, 8093, Zurich, Switzerland;" ]
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[ "astronomy" ]
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[ "instabilities", "planet–disk interactions", "planets and satellites: formation", "protoplanetary disks", "Astrophysics - Earth and Planetary Astrophysics" ]
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[ "10.3847/0004-637X/832/2/178", "10.48550/arXiv.1607.05775" ]
1607
1607.05775_arXiv.txt
Protoplanetary disks are the cradles of young planets, where several dynamical processes are likely involved in the planet formation process \citep[e.g.][]{Armitage2011}. Of particular interest are the transition disks, disks with inner millimeter-dust cavities. In the last decade, observations have revealed that {some} transition disks are far from axisymmetric: azimuthal asymmetries in the submillimeter continuum are thought to be dust traps, triggered by vortices acting as azimuthal pressure bumps \citep[e.g.][]{vanderMarel2013,LyraLin2013,Birnstiel2013}. On the other hand, near-infrared scattered light observations show large spirals \citep[e.g.][]{Muto2012,Garufi2013,Grady2013,Avenhaus2014}. Both spirals and vortices {may} indicate the presence of recently formed massive planets: in the case of a vortex through Rossby wave instability (RWI) {at the steep edges} of the gap that is carved by the planet \citep[][]{Lovelace1999,deVal-Borro2007} and in the case of spirals through the trigger of density waves directly by the planet {\citep[e.g.][]{Kley2012}}. \begin{figure*}[!ht] \begin{center} \epsscale{2} \plotone{ALMA_basic.pdf} \end{center} \caption{{\bf a.} 335 GHz continuum emission of HD~135344B in superuniform weighting. {\bf b.} Overlay of the PDI image of \citet{Garufi2013} (black contours) on top of the ALMA continuum emission. The spirals as identified by \citet{Muto2012} are labeled as S1 and S2. {\bf c.} PDI image of \citet{Garufi2013} in blue colors. In a and c, the white dashed ellipse indicates the 45 AU radius. \label{fig:imagedata}} \end{figure*} Alternative explanations for spiral arms in disks include RWI at the edge of a dead zone \citep{lyra2015}, accretion from an envelope \citep{lesur2015} and gravitational instability \citep{Lodato2004,Lodato2005,Rice2004}, although estimated disk masses generally appear to be too low for them to be self-gravitating \citep{WilliamsCieza2011}. A natural question is whether there is any relation between the spiral arms observed in near-infrared scattered light (from the disk surface) and the structures seen in submillimeter emission (from the midplane). Although spiral features in submillimeter emission have been seen in two transition disks \citep{Pietu2005,Christiaens2014}, they are not entirely consistent with their near infrared counterparts. \citet{Juhasz2015,Pohl2015,Dong2015spirals} demonstrated that spirals generated by planet-disk interaction more likely results from changes in the vertical structure rather than the density structure, which are hard to detect in millimeter emission. On the other hand, spirals that form through gravitational instability can trap dust \citep{Lodato2004,DiPierro2015,Dong2015GI}, resulting in millimeter continuum spirals. In this Letter we present Atacama Large Millimeter/submillimeter Array (ALMA) submillimeter continuum observations at very high spatial resolution of HD~135344B \footnote{also known as SAO~206462} (F4 star, $d\sim$140 pc, $\sim$8 Myr \citep{vanBoekel2005,Grady2009}), a well-studied transition disk at both optical and millimeter wavelengths. The HD~135344B disk contains a $\sim$40 AU radius dust cavity \citep{Brown2007,Brown2009,Andrews2011} with a minor azimuthal asymmetry along the dust ring \citep{Perez2014,Pinilla2015beta}. CO observations and scattered light indicate that gas and small grains are present inside the cavity \citep{Pontoppidan2008,Lyo2011,vanderMarel2015-12co,vanderMarel2015-isot,Garufi2013}, consistent with a scenario where a massive planet at $\lesssim$30 AU has cleared its orbit and trapped the large dust further out \citep{Pinilla2012b}. Scattered light imaging reveal two major spiral arms \citep{Muto2012, Garufi2013, Stolker2016}, proposed to be linked to planet-disk interaction, with planets located at 55 and 126 AU. The new images presented in this letter show substructure in the millimeter emission to an unprecedented level, revealing a double structure, which may be responsible for triggering the spiral arms seen in the scattered light. This new interpretation has consequences for the implied location of the putative planets.
\begin{figure*}[!ht] \epsscale{2} \plotone{Newcartoon.pdf} \caption{Cartoon explaining the proposed scenario.} \label{fig:cartoon} \end{figure*} The F1 feature is not consistent with the spiral arm prescription, but it can be described as a ring ($\sim50$ AU) with an asymmetry at $\sim80$ AU. Therefore we propose a new alternative scenario for this disk to explain the structure of both millimeter and scattered light data. The millimeter geometry is consistent with a model from \cite{Lobo-Gomes2015}, showing that a planet generates a pressure bump at $50$ AU (F2), which triggers a second generation vortex at $80$ AU (F1). The cavity radius of the gas and small grains \citep{Garufi2013,vanderMarel2015-isot} suggests the presence of a massive planet at 30 AU. A millimeter dust ring at 50 AU (F2) is consistent with this scenario, as the dust is trapped further out than the gas gap edge \citep{Pinilla2012b}. The ALMA and PDI data trace different grain size populations and disk heights, possibly driven by different mechanisms. However, it is striking that F1 coincides with the edge of the S1 arm. We propose that the S1 is triggered by a vortex that has created the dust asymmetry F1, since vortices can be massive enough to launch their own density waves in a disk when self-gravity is included in hydrodynamical models \citep[e.g.][]{BaruteauZhu2015}. Only a lower limit to the mass of the F1 feature can be set as the emission is partially optically thick: with a total flux of $\sim$200 mJy and a ISM gas-to-dust ratio of 100, the total mass is $>$16 M$_{\rm Jup}$ (using $M_{\rm gas} = 0.08*F_{\nu}(d/140 {\rm pc})^2$ M$_{\rm Jup}$, \citet{Cieza2008}). The outer extent of S1 (outside the vortex) remains undetectable in the PDI image due to the lower brightness in the outer disk. \citet{Muto2012} find a best-fit for the launching point of S1 at $r_c$=0.39" (55 AU) at $\theta_0$=204$^{\circ}$, but with a large confidence interval (see Figure 5 in Muto et al.). Fitting the S1 spiral with an initial guess close to the center of the vortex results in the fit in Figure \ref{fig:fitS12}a with $r_c,\theta_0=0.6'',180^{\circ}$ (84 AU) and $h_c=0.08$. This launching point does not coincide exactly with the center of F1, although there is a large uncertainty due to the unknown scale height at this location. Furthermore, ALMA continuum observations trace the mm-dust, whose center may not coincide with the gas vortex \citep{BaruteauZhu2015}, and the vortex can be a large scale structure where the center of mass may not be well represented by a single location, contrary to a planet. On the other hand, the S2 spiral was best-fit by \citet{Muto2012} for $r_c,\theta_0$=0.9" (126 AU), 353$^{\circ}$, but we find that it can also be fit with a launching point in the inner part of the disk for $r_c,\theta_0$=0.23" (32 AU), 211$^{\circ}$ (Figure \ref{fig:fitS12}b). The launching point of S2 would be a massive planet, just inside the gas cavity radius \citep{vanderMarel2015-isot}. \citet{Stolker2016} finds a best fit for the S2 launching point to the VLT/SPHERE data slightly further in, at $r_c,\theta_0$=0.15" (21 AU), 247$^{\circ}$. We propose that the combination of the scattered light and the millimeter observations is consistent with the following sequence of events (see Figure \ref{fig:cartoon}): \begin{enumerate} \item A massive planet is formed at $\sim$30 AU radius. \item The planet triggers a spiral density wave outwards (PDI S2 feature). \item The planet clears its orbit in the gas (CO observations) and creates a radial pressure bump at its edge where millimeter-dust gets trapped (ALMA continuum F2 feature). \item The pressure bump creates an effective $\alpha$ viscosity that is large enough to induce accretion, depleting the gas and inducing a second pressure bump further out. The second pressure bump triggers RWI, forms a vortex and traps the millimeter-dust asymmetrically (ALMA continuum F1 feature). \item The outer vortex triggers a spiral density wave inwards (PDI S1 feature). \end{enumerate} This scenario can potentially explain both PDI and millimeter observations. Hydrodynamical models of gas and dust, including self-gravity, are required to check whether our proposed scenario can instead quantitatively explain the observed structures of HD135344B. One of the major uncertainties in the scenario are the fits to the locations of the launching points. The reason is that the scattered light data are mainly sensitive to changes in the scale height and therefore, the observed scattered light is significantly affected by geometric parameters. The observed spirals form only the illuminated inner part of a surface change. Also, the inner disk region may shadow the outer part and thus alter the intrinsic disk scale height distribution. In particular, the azimuthal angle of the continuum ALMA feature coincides with the brighter part of the closer-in S2 spiral and therefore, S2 may be casting a shadow on part of S1, affecting the fit of the launching points. Another caveat is the symmetry of the two spiral arms at the time of observation, suggesting a common nature such as proposed by \citet{Dong2015spirals} who demonstrates the trigger of two symmetric spiral arms by a single planet at 100 AU. As this planet has remained undetected, this scenario cannot be confirmed. On the other hand, if there are instead two launching points (32 and 86 AU), the two spirals would have distinct angular velocities and their symmetric appearance is fortuitous, making the scenario less probable. The orbital period of the 32 AU point is only 143 years, implying a 2.5$^{\circ}$/year angular shift. Repeating the scattered light observations in 5 years should clearly reveal the motion of this arm. If the asymmetry is indeed related to a vortex, an azimuthal shift of $\sim$0.1" (6$^{\circ}$) in the millimeter continuum (measurable at 0.2" resolution) is detectable after 10 years. The scenario is an example of triggered planet formation, where the formation of a first planet can induce dust growth and potentially further planet formation in the outer disk.
16
7
1607.05775
In recent years, spiral structures have been seen in scattered light observations and signs of vortices in millimeter images of protoplanetary disks, both probably linked with the presence of planets. We present Atacama Large Millimeter/submillimeter Array Band 7 (335 GHz or 0.89 mm) continuum observations of the transition disk HD 135344B at an unprecedented spatial resolution of 0.″16, using superuniform weighting. The data show that the asymmetric millimeter-dust ring seen in previous work actually consists of an inner ring and an outer asymmetric structure. The outer feature is cospatial with the end of one of the spiral arms seen in scattered light, but the feature itself is not consistent with a spiral arm due to its coradiance. We propose a new possible scenario to explain the observed structures at both wavelengths. Hydrodynamical simulations show that a massive planet can generate a primary vortex (which dissipates at longer timescales, becoming an axisymmetric ring) and trigger the formation of a second generation vortex further out. Within this scenario, the two spiral arms observed at scattered light originate from a planet at ∼30 au and from the secondary vortex at ∼75 au rather than a planet further out as previously reported.
false
[ "vortices", "planets", "superuniform weighting", "protoplanetary disks", "longer timescales", "spiral structures", "scattered light observations", "millimeter images", "∼30 au", "a second generation vortex", "continuum observations", "scattered light", "HD 135344B", "the secondary vortex", "a primary vortex", "previous work", "second", "a massive planet", "Array Band", "an axisymmetric ring" ]
9.560342
13.006043
149
1867405
[ "Zink, Jonathon K.", "Christian, Damian" ]
2016arXiv160707894Z
[ "The search for IR excess in low signal to noise sources" ]
0
[ "-", "-" ]
null
[ "astronomy" ]
3
[ "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
null
[ "10.48550/arXiv.1607.07894" ]
1607
1607.07894_arXiv.txt
Circumstellar disks are created from the remnant material of stellar formation. Young protostellar disks provide a method for distinguishing the age of its stellar host and help models converge in determining the exact mechanisms of planet formation. Current models suggest most protoplanetary disks will photoevaporate within $\sim$ 1-5 Myr (Alexander et al 2006a, b; Owen et al. 2010). This infancy in which the star has just begun fusion, but not yet shed its disk, is the key time in which planet formation occurs. Finding stars within this narrow window of the stars lifetime, provides a further glimpse into the mysterious cause of planet formation. Additional clues to planet formation have resulted from the many planetary systems with large dust disks (Kalas et al. 2008; Lagrange et al. 2010; Marois et al. 2008; 2010). The presence of holes, gaps, and azimuthal symmetries in the dust distribution may also indicate the presence of undiscovered planets. Although many studies have not found strong correlation between the presence of circumstellar disks and planets, newer \textit{Herschel} observations have suggested there is a correlation (Marshall et al. 2014; Kennedy et al. 2014; 2015). For an alternate view see Moro-Martín et al. 2015. \par There have been many studies attempting to quantify the occurrence of IR excesses and their inferred disks in FKG and M type stars. The occurrence of excess IR emission at longer wavelengths (70 $\mu$m), than those found by the mid IR régime of this study, have been found to be 10-15\% (Beichman et al. 2006; Trilling et al. 2008), compared to a much lower rate of $\sim$1\% for 24 $\mu$m emission (Lawler et al. 2009). Expanding these samples to stars known to host planets has found similar or even slightly lower rate for the occurrence of IR excesses (Bryden et al. 2009). More recently, the Wide-field Infrared Survey Explorer (WISE) provides information on millions of stars at 22 $\mu$m and Morales et al. (2012) found nine planet-bearing stars with warm dust emission; this gives an excess incidence for planet-bearing of only 1\% for Main Sequence stars. \par Here we have undertaken a study to select stars that provide evidence of a disk from the ALLWISE catalog. This study differs from Patel et al. (2014), who searched for dust disks in the \textit{Hipparcos} catalog utilizing WISE data, Avenhaus et al. (2012), who detected circumstellar disks around M dwarfs, and the Theissen et al. (2014) study, which sought to examine population synthesis among disk harboring stars, by focusing on low SNR sources (<15) and further accounts for reddening effects seen by high magnitude signals in the WISE database. We also re-examine the current list of Kepler candidates for possible excess candidates (initial study was performed by Kennedy and Wyatt 2012, known as KW12 from here forth). In Section 2, we present the target selection criteria and the WISE photometry used. In Section 3 we present the IR excess results, and a table of their important parameters. In Section 4 we investigate some of the key features of the candidates, present Spectral Energy Distributions (SEDs) for noteworthy sources. Finally, in Section 5 concluding remarks are provided.
This study presents 14 candidates that show statically significant IR excess near the sensitive limits of WISE using this weighted method. These are likely due to dust disks surrounding host stars. Each candidate has been thoroughly vetted to ensure true excess over erroneous possibilities. Further combing of the Kepler candidates, found one star with disk like features that provide robust excess over background contamination. This star warrants further examination by direct imaging follow up. Of the systems discovered, one presents a large disk orbit 8.40 $\pm$ 0.73 AU and makes a good target for future deep imaging with AO and coronagraphs to search for new exoplanets. Several other stars provide lower bounds on disk radius and may harbor much larger disk, capable of imaging detection. This study provides evidence that true disk sources are abundant in these low SNR regions and merit continued study. \begin{figure} \resizebox{\hsize}{!}{\includegraphics{fig7.pdf}} \caption{Four WISE waveband images for TYC 3143-322-1 (Top) and J053010.20-010140.9 (Bottom) . Represented from left to right are W1-4 bands respectively. These images show the mild contamination in these bands, indicating a true point source. A \SI{75}{\arcsecond} x \SI{75}{\arcsecond} portion of the sky is displayed in each image. } \end{figure}
16
7
1607.07894
We present sources selected from their Wide-field Infrared Survey Explorer (WISE) colors that merit future observations to image for disks and possible exoplanet companions. Introducing a weighted detection method, we eliminated the enormous number of specious excess seen in low signal to noise objects by requiring greater excess for fainter stars. This is achieved by sorting through the 747 million sources of the ALLWISE database. In examining these dim stars, it can be shown that a non-Gaussian distribution best describes the spread around the main-sequence polynomial fit function. Using a gamma Probability Density Function (PDF), we can best mimic the main sequence distribution and exclude natural fluctuations in IR excess. With this new methodology we re-discover 25 IR excesses and present 14 new candidates. One source (J053010.20-010140.9), suggests a 8.40 $\pm$ 0.73 AU disk, a likely candidate for possible direct imagining of planets that are likely fully formed. Although all of these sources are well within the current flux ratio limit of $\sim$10$^{-6}$ (Wyatt 2008), J223423.85+403515.8 shows the highest bolometric flux ratio ($f_d$=0.0694) between disk and host star, providing a very good candidate for direct imaging of the circumstellar disk itself. In re-examining the Kepler candidate catalog (original study preformed by Kennedy and Wyatt 2012), we found one new candidate that indicates disk like characteristics (TYC 3143-322-1).
false
[ "IR excess", "disk", "disks", "greater excess", "possible exoplanet companions", "possible direct imagining", "fainter stars", "specious excess", "direct imaging", "non-Gaussian", "disk and host star", "noise objects", "low signal", "TYC", "future observations", "natural fluctuations", "IR", "Infrared Survey Explorer", "Wyatt", "original study" ]
9.00595
12.922551
-1
12435116
[ "Henrot-Versillé, S.", "Perdereau, O.", "Plaszczynski, S.", "Rouillé d'Orfeuil, B.", "Spinelli, M.", "Tristram, M." ]
2016arXiv160702964H
[ "Agnostic cosmology in the CAMEL framework" ]
8
[ "-", "-", "-", "-", "-", "-" ]
[ "2017A&A...606A.104C", "2019A&A...623A...9H", "2019MNRAS.482.4290H", "2019PDU....24..260B", "2021MNRAS.504.5840F", "2023PDU....4201348P", "2023PhRvD.108l3514H", "2024arXiv240504455G" ]
[ "astronomy" ]
5
[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1953JChPh..21.1087M", "1975CoPhC..10..343J", "1980CoPhC..20...29J", "1992StaSc...7..457G", "1996NIMPA.368..793P", "1998PhRvD..57.3873F", "2001CQGra..18.2677C", "2002PhRvD..66j3511L", "2004AIPC..735..395S", "2005MNRAS.356..925D", "2005MNRAS.358..833T", "2007ApJ...654....2F", "2009MNRAS.398.1601F", "2010bmic.book.....H", "2011MNRAS.415.2892B", "2011MNRAS.416.3017B", "2012ApJ...755...70R", "2012MNRAS.427.2132P", "2012MNRAS.427.3435A", "2013ApJ...779...86S", "2013JCAP...02..001A", "2013PASP..125..306F", "2014A&A...566A..54P", "2014A&A...568A..22B", "2014MNRAS.437.3918A", "2014MNRAS.441...24A", "2015A&C....12...45Z", "2015ApJ...799..177G", "2015CQGra..32d5003H", "2015MNRAS.450L..61H", "2016A&A...594A..11P", "2016A&A...594A..13P", "2016A&A...594A..15P", "2016A&A...596A.108P", "2017A&A...597A.126C" ]
[ "10.48550/arXiv.1607.02964" ]
1607
1607.02964_arXiv.txt
\label{sec:Introduction} Since the 2000's \citep{Christensen_2001}, the adoption of Monte Carlo Markov Chain (MCMC) techniques by the cosmological community has promoted the acceptance of Bayesian methodology for parameter estimation. The possible (although inefficient) sampling of a high dimensional space allows, invoking Bayes theorem, to reconstruct posterior distributions assuming some prior degree of belief (rarely discussed). MCMC usage was popularized in cosmology by the \COSMOMC\ package \citep{cosmomc} which uses an optimized version of the Metroplis-Hastings algorithm (see also \texttt{MontePython} \citep{montepython}). There exist today several implementations of other algorithms, as \texttt{MultiNest} \citep{multinest} or \texttt{PolyChord} \citep{Handley_2015} for Nested Sampling \citep{skilling2004}, or \texttt{emcee} \citep{emcee} for the affine invariant ensemble sampler \citep{goodman}. Several of them are packaged within the \texttt{CosmosSIS} package \citep{cosmosis} and some of them were compared in \citet{Allison_2014}. They are essentially written in \texttt{Fortran}/\texttt{python} and/or wrapped into multiple languages. All of them adopt the Bayesian paradigm (see \citet{BayesCosmo} for a review). In other communities, as High Energy Physics, the frequentist approach is more traditional. The reason for its absence in cosmology maybe lies in the difficulty of building precise profile-likelihoods because of the numerical noise inherent to the Boltzmann solver computations (Sect.\ref{sec:precision}). It was however shown to be feasible on real \planck\ data using an accurate minimizer and tuning the precision of the Boltzmann computations \citep{prof14}. We wish to avoid the ideological debate about "who is right", and focus on the interest of confronting both methodologies. It is the goal of the \camel package to provide tools to compare the Bayesian and frequentist analyzes in order to better understand for instance the Bayesian priors effects or the shape of the likelihood. We choose to write it in pure \texttt{C++} since we consider it as a proper language to develop a robust, clean and long term project. By \texttt{C++} we mean an object-oriented code with some level of abstraction and well defined design patterns in order to build a clean interface that anticipates future developments. Fortunately users do not need to know what is under the hood and one can plug in easily any new Boltzmann solver and/or likelihood. Certainly the best way to present \CAMEL\ is by working out a full use-case. We will revisit the \lcdm\ cosmological parameters estimation with \planck\ data using the \hlp\ likelihood \citep{Like15} and compare the Bayesian and frequentist results. Not only will we show how to obtain precise best-fits and profile-likelihoods, but also propose a new implementation of an Adaptive scheme for the Metropolis MCMC algorithm that relieves the pain of reconstructing the proposal. Sect. \ref{sec:pipeline} is a quick overview of the building blocks of parameter estimation. Then Sect. \ref{sec:use_case} explains in minute-details how to obtain best-fits (Sect.\ref{sec:MLE}), profile-likelihoods (Sect.\ref{sec:profile}) and produce MCMC with the new Adaptive algorithm (Sect. \ref{sec:MCMC}). For each method, the basics will be reviewed. This part concretely shows how to produce all the results and can therefore be used as a tutorial. Sect \ref{sec:comp} finally shows the interest of comparing both methods focusing on the study of the likelihood "volume effects". These results, partially available in \citet{prof14} and \citet{cosmo15}, were however never compiled and detailed. The results on the volume effects that affect \planck's \hiell\ likelihoods are new and complement \citet{papierAlens}.
We have seen how comparing different statistical approaches can be fruitful in cosmological analyzes. A possible workflow can be: \begin{enumerate} \item first run an accurate Minimization with \minuit with a few starting points. Obtain the best-fit and the Hessian matrix. \item use the Hessian as the starting matrix to the proposal of the Adaptive Metroplis MCMC algorithm. \item if some posterior modes differ significantly from the best-fit value, run single profile-likelihoods on those parameters. \item if some parameters are of particular importance (as $(w_0,w_a)$ in future surveys \footnote{2D profile-likelihoods are also implemented}), always run a profile-likelihood on them to verify the influence of priors and volume effects in the MCMC run. \end{enumerate} \camel provides the tools to perform such rich studies \textit{using the very same input file}. See the \url{camel.inp3.fr} web site for more information. The software is collaborative and contributions are most welcome.
16
7
1607.02964
Cosmological parameter estimation is traditionally performed in the Bayesian context. By adopting an "agnostic" statistical point of view, we show the interest of confronting the Bayesian results to a frequentist approach based on profile-likelihoods. To this purpose, we have developed the Cosmological Analysis with a Minuit Exploration of the Likelihood ("CAMEL") software. Written from scratch in pure C++, emphasis was put in building a clean and carefully-designed project where new data and/or cosmological computations can be easily included. CAMEL incorporates the latest cosmological likelihoods and gives access from the very same input file to several estimation methods: (i) A high quality Maximum Likelihood Estimate (a.k.a "best fit") using MINUIT ; (ii) profile likelihoods, (iii) a new implementation of an Adaptive Metropolis MCMC algorithm that relieves the burden of reconstructing the proposal distribution. We present here those various statistical techniques and roll out a full use-case that can then used as a tutorial. We revisit the $\Lambda$CDM parameters determination with the latest Planck data and give results with both methodologies. Furthermore, by comparing the Bayesian and frequentist approaches, we discuss a "likelihood volume effect" that affects the optical reionization depth when analyzing the high multipoles part of the Planck data. The software, used in several Planck data analyzes, is available from http://camel.in2p3.fr. Using it does not require advanced C++ skills.
false
[ "new data", "several Planck data analyzes", "several estimation methods", "Maximum Likelihood Estimate", "Adaptive Metropolis MCMC", "Cosmological parameter estimation", "cosmological computations", "http://camel.in2p3.fr", "advanced C++ skills", "A high quality Maximum Likelihood Estimate", "results", "the latest Planck data", "the latest cosmological likelihoods", "Bayesian", "the Planck data", "an Adaptive Metropolis MCMC algorithm", "profile-likelihoods", "iii", "pure C++", "part" ]
11.571186
2.985587
-1
2495381
[ "Fioretti, Valentina", "Bulgarelli, Andrea", "Malaguti, Giuseppe", "Spiga, Daniele", "Tiengo, Andrea" ]
2016SPIE.9905E..6WF
[ "Monte Carlo simulations of soft proton flares: testing the physics with XMM-Newton" ]
12
[ "INAF Istituto di Astrofisica Spaziale e Fisica Cosmica Bologna (Italy)", "INAF Istituto di Astrofisica Spaziale e Fisica Cosmica Bologna (Italy)", "INAF Istituto di Astrofisica Spaziale e Fisica Cosmica Bologna (Italy)", "INAF Osservatorio Astronomico di Brera (Italy)", "INAF Istituto di Astrofisica Spaziale e Fisica Cosmica Milano (Italy)" ]
[ "2016SPIE.9905E..5OS", "2017ExA....44..321G", "2018ApJ...866..126H", "2018SPIE10699E..3JF", "2020Ap&SS.365..158Z", "2020ApJ...903...89K", "2020SPIE11444E..2OW", "2021JGRA..12629273K", "2022APh...13702668Z", "2022ApJ...928..168G", "2023A&A...672A.156O", "2023ApJ...947...49H" ]
[ "astronomy", "physics" ]
3
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
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[ "10.1117/12.2232537", "10.48550/arXiv.1607.05319" ]
1607
1607.05319_arXiv.txt
\label{sec:intro} % The Earth magnetic field confines the interplanetary energetic plasma (mainly protons and electrons) in two distinct radiation belts (the Van Allen belts). The inner one, more compact and mostly filled with protons, extends from about 700 to 10$^{4}$ km. The outer one, recently discovered to occasionally split into a third radiation ring\cite{2013Sci...340..186B}, is mainly constituted by high energy electrons and stretches up to $6\times10^{4}$ km. The trapped particles can pose a significant radiation threat to the on-board electronic systems, depending on the telescope orbit and technology. Launched in July and December 1999 respectively, the NASA Chandra and the ESA XMM-Newton X-ray missions represented a cornerstone in the X-ray astronomical exploration, thanks to their unprecedented sensitivity and angular resolution below 10-15 keV. They are currently operating in a high eccentricity elliptical orbit that, reaching extreme apogees ($>10^{5}$ km), allows them to spend most of the operating time beyond the radiation belts and without the frequent interruption of the Earth's shadowing. On the contrary, in the perigee passage the spacecraft altitude is in the range $(5-10)\times10^{3}$ km, forcing the telescopes to cross the radiation belts each orbit. \\ Both telescopes carry Wolter type-I mirrors to focus X-rays through grazing angle reflection to the detection plane, which is shielded by heavy absorbers outside the field of view to block penetrating radiation (mainly cosmic rays) to reach the focal plane. Electrons in the radiation belts and in the outer regions can also be scattered by the X-ray optics and funnelled to the focal plane. For this reason, X-ray telescopes on board Chandra, XMM-Newton and Swift are equipped with magnetic diverters that deflect their path outside the detection plane (see e.g Ref. \citenum{wil00}). What came unexpected after one month of operation of the Chandra mission was that also low energy protons entering the field of view at grazing angles can be reflected too by the X-ray mirror shells and reach the focal plane. With energies below 100-300 keV, these so-called "soft protons" can cause serious damages to the electronic system and the overall mission performance. \\ XMM-Newton's filter wheel completely blocks the European Photon Imaging Camera (EPIC) field of view when crossing the radiation belts protecting the detectors from damage. However, the soft protons populating the outer magnetosphere, although with a $\sim10^3$ lower flux\cite{2000SPIE.4140...99O}, including the magnetotail, the magnetosheath and the solar wind, both in the form of a steady flux and violent Coronal Mass Ejections (CMEs), can produce photon-like deposits in the read-out system that increase the X-ray residual background level and even threaten the observation itself \cite{bri00}. The performance of future X-ray focusing telescopes operating outside the radiation belts (e.g., the ESA Athena\cite{2013arXiv1306.2307N} mission and the eROSITA\cite{2014SPIE.9144E..1TP} instrument on-board the Russian Spektr-RG observatory, both to be placed in L2 orbit) might also be affected by soft proton contaminations. Monte Carlo simulations are mandatory to evaluate the impact of such events and accordingly design shielding solutions (e.g., a magnetic diverter) without limiting the sensitivity of the instruments, and require the accurate characterization of (i) the distribution of the soft proton population, (ii) the mirror-proton physics interaction at play, and (iii) the effect on the focal plane. Despite many solutions proposed so far to explain the physics interaction behind the soft proton grazing angle scattering, the difficulty in obtaining accurate laboratory measurements\cite{2015ExA....39..343D} prevents ray-tracing codes from implementing physically-reliable models. After presenting an overview of the XMM-Newton soft proton flare behaviour (Sec. \ref{sec:obs}), we test the current knowledge of the input proton population and the physics interaction (Sec. \ref{sec:physics}) by simulating the soft proton interaction with a simplified model of the X-ray mirror module and the EPIC/pn camera. The final aim is comparing for the first time the simulated X-ray spectrum with a real on-flight XMM-Newton observation of a soft proton flare (Sec. \ref{sec:sim}). In the present work, we use Geant4 9.1 - current release at the time of writing being 10.2 - to cross-check our results with the ESA Geant4 simulation campaign (see e.g. [\citenum{nar02a}] outcome after the launch of Chandra.
The Geant4-based BoGEMMS framework has been used to simulate the EPIC/pn X-ray background flux induced by soft proton scattering on XMM-Newton optics. The Equator-S proton flux, detected at an altitude of 70000 km in the 30-1000 keV energy range, has been used as input proton population, and the combination of Firsov (at proton incident angles $<1$ deg.) and mutiple (at proton incident angles $>1$ deg.) scattering models has been applied to model the interaction with the X-ray shells. The resulting background flux at 1 keV is $6\times10^{-3}$ counts cm$^{-2}$ s$^{-1}$ keV$^{-1}$. \\ When comparing the simulation with the real X-ray spectrum detected by the pn and the MOS instruments of the EPIC camera in two intense soft proton flare episodes, our result underestimates up to two orders of magnitude the real case. However, it must be noted that soft proton flares are extremely variable, so that one single observation is not representive of the general behaviour. The solution would be the use of a time averaged input proton flux, and an average proton spectrum obtained by folding together soft proton flares detected along several orbits. The present results must be intended as a preliminary step in the comprehension of low energy proton scattering by X-ray optics. Simulations with the most recent Geant4 releases, providing more accurate models for the multiple scattering, and the implementation of new models for proton grazing angle scattering on the basis of laboratory measurements and ray-tracing simulations, are planned in the near future. The final aim is not only reproducing XMM-Newton soft proton flares, but more importantly the correct evaluation of this important background component for the design of future X-ray space telescopes (e.g the ESA ATHENA mission).
16
7
1607.05319
Low energy protons (&lt; 100 - 300 keV) in the Van Allen belt and the outer regions can enter the field of view of X-ray focusing telescopes, interact with the Wolter-I optics, and reach the focal plane. The funneling of soft protons was discovered after the damaging of the Chandra/ACIS Front-Illuminated CCDs in September 1999 after the first passages through the radiation belt. The use of special filters protects the XMM-Newton focal plane below an altitude of 70000 km, but above this limit the effect of soft protons is still present in the form of sudden ares in the count rate of the EPIC instruments that can last from hundreds of seconds to hours and can hardly be disentangled from X-ray photons, causing the loss of large amounts of observing time. The accurate characterization of (i) the distribution of the soft proton population, (ii) the physics interaction at play, and (iii) the effect on the focal plane, are mandatory to evaluate the background and design the proton magnetic diverter on board future X-ray focusing telescopes (e.g. ATHENA). Several solutions have been proposed so far for the primary population and the physics interaction, however the difficulty in precise angle and energy measurements in laboratory makes the smoking gun still unclear. Since the only real data available is the XMM-Newton spectrum of soft proton flares in orbit, we try to characterize the input proton population and the physics interaction by simulating, using the BoGEMMS framework, the proton interaction with a simplified model of the X-ray mirror module and the focal plane, and comparing the result with a real observation. The analysis of ten orbits of observations of the EPIC/pn instrument show that the detection of flares in regions far outside the radiation belt is largely influenced by the different orientation of the Earth's magnetosphere respect with XMM-Newton'os orbit, confirming the solar origin of the soft proton population. The Equator-S proton spectrum at 70000 km altitude is used for the proton population entering the optics, where a combined multiple and Firsov scattering is used as physics interaction. If the thick filter is used, the soft protons in the 30-70 keV energy range are the main contributors to the simulated spectrum below 10 keV. We are able to reproduce the proton vignetting observed in real data-sets, with a 50% decrease from the inner to the outer region, but a maximum flux of 0:01 counts cm<SUP>2</SUP> s<SUP>-1</SUP> keV<SUP>-1</SUP> is obtained below 10 keV, about 5 times lower than the EPIC/MOS detection and 100 times lower than the EPIC/pn one. Given the high variability of the are intensity, we conclude that an average spectrum, based on the analysis of a full season of soft proton events is required to compare Monte Carlo simulations with real events.
false
[ "soft proton events", "soft protons", "soft proton flares", "X-ray focusing telescopes", "X-ray photons", "board future X-ray focusing telescopes", "Low energy protons", "physics interaction", "real events", "observing time", "EPIC", "the X-ray mirror module", "the soft proton population", "pn instrument", "regions", "Monte Carlo simulations", "the input proton population", "the soft protons", "energy measurements", "Monte Carlo" ]
9.779549
3.432116
67
12472677
[ "de la Fuente Marcos, C.", "de la Fuente Marcos, R." ]
2016MNRAS.462.1972D
[ "Finding Planet Nine: apsidal anti-alignment Monte Carlo results" ]
21
[ "Apartado de Correos 3413, E-28080 Madrid, Spain", "Apartado de Correos 3413, E-28080 Madrid, Spain" ]
[ "2016PhDT.......541S", "2017AJ....153...63S", "2017AJ....153...65M", "2017Ap&SS.362...11I", "2017Ap&SS.362..198D", "2017MNRAS.467L..66D", "2017MNRAS.471L..61D", "2017RNAAS...1....5D", "2018AJ....155..166M", "2018AJ....156...81B", "2018AJ....156..263L", "2018MNRAS.474..838D", "2018RNAAS...2..167D", "2018exha.book.....P", "2019PhDT........68B", "2019PhR...805....1B", "2020tnss.book.....P", "2020tnss.book...61K", "2020tnss.book...79T", "2021MNRAS.506..633D", "2024arXiv240115808D" ]
[ "astronomy" ]
6
[ "methods: statistical", "celestial mechanics", "minor planets", "asteroids: general", "Oort Cloud", "planets and satellites: detection", "planets and satellites: general", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1996DPS....28.2504G", "1999ssd..book.....M", "2014ApJ...781....4L", "2014MNRAS.443L..59D", "2014Natur.507..471T", "2015ApJ...806...42K", "2015Icar..258...37G", "2015MNRAS.446.1867D", "2015MNRAS.446.3788B", "2015MNRAS.453.3157J", "2015MNRAS.454.3267F", "2016A&A...587L...8F", "2016A&A...589A..63P", "2016A&A...589A.134L", "2016A&A...590L...2B", "2016AJ....151...22B", "2016ApJ...822L...2C", "2016ApJ...822L..11G", "2016ApJ...823L...3L", "2016ApJ...824L..22M", "2016ApJ...824L..23B", "2016ApJ...824L..25F", "2016ApJ...825...33K", "2016ApJ...826...64B", "2016MNRAS.456.1587B", "2016MNRAS.457L..89M", "2016MNRAS.459L..66D", "2016MNRAS.460L..64D", "2016MNRAS.460L.109M", "2016MNRAS.460L.123D" ]
[ "10.1093/mnras/stw1778", "10.48550/arXiv.1607.05633" ]
1607
1607.05633_arXiv.txt
The distribution of the orbital parameters of the known extreme trans-Neptunian objects or ETNOs is statistically incompatible with that of an unperturbed asteroid population following Keplerian orbits (de la Fuente Marcos \& de la Fuente Marcos 2014, 2016b; Trujillo \& Sheppard 2014; de la Fuente Marcos, de la Fuente Marcos \& Aarseth 2015, 2016; Gomes, Soares \& Brasser 2015; Batygin \& Brown 2016; Brown \& Batygin 2016; Malhotra, Volk \& Wang 2016). A number of plausible explanations have been suggested. These include the possible existence of one (Trujillo \& Sheppard 2014; Gomes et al. 2015; Batygin \& Brown 2016; Brown \& Batygin 2016; Malhotra et al. 2016) or more (de la Fuente Marcos \& de la Fuente Marcos 2014, 2016b; de la Fuente Marcos et al. 2015, 2016) trans-Plutonian planets, capture of ETNOs within the Sun's natal open cluster (J\'{\i}lkov\'a et al. 2015), stellar encounters (Brasser \& Schwamb 2015; Feng \& Bailer-Jones 2015), being a by-product of Neptune's migration (Brown \& Firth 2016) or the result of the inclination instability (Madigan \& McCourt 2016), and having been induced by Milgromian dynamics (Pau\v{c}o \& Kla\v{c}ka 2016). At present, most if not all of the unexpected orbital patterns found for the known ETNOs seem to be compatible with the trans-Plutonian planets paradigm that predicts the presence of one or more planetary bodies well beyond Pluto. Within this paradigm, the best studied theoretical framework is that of the so-called Planet Nine hypothesis, originally suggested by Batygin \& Brown (2016) and further developed in Brown \& Batygin (2016). The goal of this analytically and numerically supported conjecture is not only to explain the observed clustering in physical space of the perihelia and the positions of the orbital poles of seven ETNOs (see Appendix A for further discussion), but also to account for other, previously puzzling, pieces of observational evidence like the existence of low perihelion objects moving in nearly perpendicular orbits. The Planet Nine hypothesis is compatible with existing data (Cowan, Holder \& Kaib 2016; Fienga et al. 2016; Fortney et al. 2016; Ginzburg, Sari \& Loeb 2016; Linder \& Mordasini 2016) but, if Planet Nine exists, it cannot be too massive or bright to have escaped detection during the last two decades of surveys and astrometric studies (Luhman 2014; Cowan et al. 2016; Fienga et al. 2016; Fortney et al. 2016; Ginzburg et al. 2016; Linder \& Mordasini 2016). A super-Earth in the sub-Neptunian mass range is most likely and such planet may have been scattered out of the region of the Jovian planets early in the history of the Solar system (Bromley \& Kenyon 2016) or even captured from another planetary system (Li \& Adams 2016; Mustill, Raymond \& Davies 2016); super-Earths may also form at 125--750 au from the Sun (Kenyon \& Bromley 2015, 2016). The analysis of the visibility of Planet Nine presented in de la Fuente Marcos \& de la Fuente Marcos (2016a) revealed probable locations of this putative planet based on data provided in Batygin \& Brown (2016) and Fienga et al. (2016); the original data have been significantly updated in Brown \& Batygin (2016). In addition, independent calculations (de la Fuente Marcos et al. 2016) show that the apsidal anti-alignment constraint originally discussed in Batygin \& Brown (2016) plays a fundamental role on the dynamical impact of a putative Planet Nine on the orbital evolution of the known ETNOs. Here, we improve the results presented in de la Fuente Marcos \& de la Fuente Marcos (2016a) focusing on the effects of the apsidal anti-alignment constraint. This paper is organized as follows. Section~2 presents an analysis of clustering in barycentric elements, pericentre and orbital pole positions, which is subsequently discussed. An updated evaluation of the visibility of Planet Nine virtual orbits at aphelion is given in Section~3. Conclusions are summarized in Section~4. \begin{table} \centering \fontsize{8}{11pt}\selectfont \tabcolsep 0.16truecm \caption{Pericentre distances, $q$, ecliptic coordinates at pericentre, $(L_q, B_q)$, and projected pole positions, $(L_{\rm p}, B_{\rm p})$, of the 16 objects discussed in this paper computed using barycentric orbits, see also Figs \ref{cluster} and \ref{poles}. (Epoch: 2457600.5, 2016 July 31.0 00:00:00.0 TDB. J2000.0 ecliptic and equinox. Input data from the SBDB; data as of 2016 July 13.) } \begin{tabular}{lccccc} \hline Object & $q$ (au) & $L_q$ (\degr) & $B_q$ (\degr) & $L_{\rm p}$ (\degr) & $B_{\rm p}$ (\degr) \\ \hline (82158) 2001 FP$_{185}$ & 34.25 & 185.28 & 3.52 & 89.36 & 59.20 \\ (90377) Sedna & 76.19 & 96.31 & $-$8.94 & 54.40 & 78.07 \\ (148209) 2000 CR$_{105}$ & 44.12 & 87.28 & $-$15.39 & 38.29 & 67.24 \\ (445473) 2010 VZ$_{98}$ & 34.35 & 71.21 & $-$3.26 & 27.40 & 85.49 \\ 2002~GB$_{32}$ & 35.34 & 213.24 & 8.49 & 87.04 & 75.81 \\ 2003~HB$_{57}$ & 38.10 & 208.32 & 2.88 & 107.87 & 74.50 \\ 2003~SS$_{422}$ & 39.42 & 359.91 & $-$8.28 & 61.05 & 73.21 \\ 2004~VN$_{112}$ & 47.32 & 35.65 & $-$13.59 & 336.02 & 64.45 \\ 2005~RH$_{52}$ & 39.00 & 336.98 & 10.83 & 216.11 & 69.55 \\ 2007~TG$_{422}$ & 35.56 & 39.41 & $-$17.88 & 22.91 & 71.40 \\ 2007~VJ$_{305}$ & 35.18 & 3.15 & $-$4.40 & 294.38 & 78.02 \\ 2010~GB$_{174}$ & 48.56 & 118.83 & $-$4.65 & 40.71 & 68.44 \\ 2012~VP$_{113}$ & 80.44 & 26.32 & $-$21.94 & 0.80 & 65.95 \\ 2013~GP$_{136}$ & 41.06 & 248.08 & 21.91 & 120.73 & 56.46 \\ 2013~RF$_{98}$ & 36.28 & 27.88 & $-$19.93 & 337.53 & 60.40 \\ 2015~SO$_{20}$ & 33.17 & 28.89 & $-$2.05 & 303.63 & 66.59 \\ \hline \end{tabular} \label{peri} \end{table} \begin{figure} \centering \includegraphics[width=\linewidth]{fperi.eps} \caption{Pericentres of the objects in Table \ref{peri}. The objects singled out in Brown \& Batygin (2016) are plotted in red. } \label{cluster} \end{figure} \begin{figure} \centering \includegraphics[width=\linewidth]{fpoles.eps} \caption{Poles of the objects in Table \ref{peri}. Those singled out in Brown \& Batygin (2016) are plotted in red, the known planets in blue, and various Planet Nine incarnations in green ---P9v0 is the nominal solution in Batygin \& Brown (2016), P9v1 is the one from Brown \& Batygin (2016), and P9 is the previous one but enforcing apsidal anti-alignment. } \label{poles} \end{figure} \begin{table*} \centering \fontsize{8}{11pt}\selectfont \tabcolsep 0.10truecm \caption{Barycentric orbital elements and parameters ---$q=a(1-e)$, $Q=a(1+e)$ is the aphelion distance, $\varpi=\Omega+\omega$, $P$ is the orbital period, $\Omega^*$ and $\omega^*$ are $\Omega$ and $\omega$ in the interval ($-\pi$, $\pi$) instead of the regular (0, 2$\pi$)--- of the known ETNOs. The statistical parameters are Q$_{1}$, first quartile, Q$_{3}$, third quartile, IQR, interquartile range, OL, lower outlier limit (Q$_{1}-1.5$IQR), and OU, upper outlier limit (Q$_{3}+1.5$IQR). Input data as in Table \ref{peri}. } \begin{tabular}{lrrrrrrrrrrr} \hline Object & $a$ (au) & $e$ & $i$ (\degr) & $\Omega$ (\degr) & $\omega$ (\degr) & $\varpi$ (\degr) & $q$ (au) & $Q$ (au) & $P$ (yr) & $\Omega^*$ (\degr) & $\omega^*$ (\degr) \\ \hline 82158 & 215.97915 & 0.84141 & 30.80134 & 179.35892 & 6.88451 & 186.24343 & 34.25244 & 397.70586 & 3172.01164 & 179.35892 & 6.88451 \\ Sedna & 506.08846 & 0.84945 & 11.92856 & 144.40251 & 311.28569 & 95.68820 & 76.19098 & 935.98594 & 11377.75735 & 144.40251 & $-$48.71431 \\ 148209 & 221.97188 & 0.80122 & 22.75598 & 128.28590 & 316.68922 & 84.97512 & 44.12267 & 399.82108 & 3304.94285 & 128.28590 & $-$43.31078 \\ 445473 & 153.36100 & 0.77602 & 4.51050 & 117.39858 & 313.72557 & 71.12415 & 34.35048 & 272.37152 & 1897.97030 & 117.39858 & $-$46.27443 \\ 2002 GB$_{32}$ & 206.50931 & 0.82887 & 14.19246 & 177.04395 & 37.04720 & 214.09115 & 35.33979 & 377.67883 & 2965.69498 & 177.04395 & 37.04720 \\ 2003 HB$_{57}$ & 159.66557 & 0.76138 & 15.50028 & 197.87107 & 10.82977 & 208.70084 & 38.09895 & 281.23218 & 2016.20147 & $-$162.12893 & 10.82977 \\ 2003 SS$_{422}$ & 197.89567 & 0.80078 & 16.78597 & 151.04690 & 209.92864 & 0.97554 & 39.42455 & 356.36679 & 2782.09180 & 151.04690 & $-$150.07136 \\ 2004 VN$_{112}$ & 327.43521 & 0.85548 & 25.54761 & 66.02280 & 326.99699 & 33.01979 & 47.32201 & 607.54841 & 5921.13675 & 66.02280 & $-$33.00301 \\ 2005 RH$_{52}$ & 153.67748 & 0.74624 & 20.44577 & 306.11067 & 32.53853 & 338.64920 & 38.99710 & 268.35786 & 1903.84838 & $-$53.88933 & 32.53853 \\ 2007 TG$_{422}$ & 502.04248 & 0.92916 & 18.59530 & 112.91071 & 285.68512 & 38.59583 & 35.56265 & 968.52230 & 11241.58933 & 112.91071 & $-$74.31488 \\ 2007 VJ$_{305}$ & 192.09934 & 0.81684 & 11.98376 & 24.38239 & 338.33491 & 2.71730 & 35.18470 & 349.01398 & 2660.76077 & 24.38239 & $-$21.66509 \\ 2010 GB$_{174}$ & 351.12735 & 0.86169 & 21.56245 & 130.71444 & 347.24510 & 117.95954 & 48.56288 & 653.69182 & 6575.27641 & 130.71444 & $-$12.75490 \\ 2012 VP$_{113}$ & 263.16564 & 0.69436 & 24.05155 & 90.80392 & 293.54965 & 24.35357 & 80.43515 & 445.89613 & 4266.39240 & 90.80392 & $-$66.45035 \\ 2013 GP$_{136}$ & 149.78673 & 0.72587 & 33.53904 & 210.72729 & 42.47818 & 253.20547 & 41.06079 & 258.51267 & 1832.00656 & $-$149.27271 & 42.47818 \\ 2013 RF$_{98}$ & 317.06525 & 0.88557 & 29.60066 & 67.53381 & 316.37528 & 23.90909 & 36.28242 & 597.84808 & 5642.08982 & 67.53381 & $-$43.62472 \\ 2015 SO$_{20}$ & 164.90289 & 0.79885 & 23.41106 & 33.63383 & 354.83023 & 28.46406 & 33.17008 & 296.63570 & 2116.21317 & 33.63383 & $-$5.16977 \\ \hline Mean & 255.17334 & 0.81082 & 20.32577 & 133.64048 & 221.52654 & 107.66702 & 43.64735 & 466.69932 & 4354.74900 & 66.14048 & $-$25.97346 \\ Std. dev. & 116.48633 & 0.06105 & 7.72495 & 71.95062 & 140.15770 & 102.40843 & 14.31074 & 226.78010 & 3103.33812 & 105.71229 & 49.06288 \\ Median & 211.24423 & 0.80903 & 21.00411 & 129.50017 & 302.41767 & 78.04963 & 38.54802 & 387.69235 & 3068.85331 & 101.85731 & $-$27.33405 \\ Q$_{1}$ & 163.59356 & 0.77236 & 15.17332 & 84.98639 & 41.12044 & 27.43644 & 35.30102 & 292.78482 & 2091.21024 & 31.32097 & $-$46.88440 \\ Q$_{3}$ & 319.65774 & 0.85096 & 24.42556 & 177.62269 & 319.26616 & 191.85778 & 44.92251 & 600.27317 & 5711.85155 & 134.13646 & 7.87082 \\ IQR & 156.06418 & 0.07860 & 9.25224 & 92.63630 & 278.14573 & 164.42135 & 9.62149 & 307.48835 & 3620.64131 & 102.81549 & 54.75523 \\ OL & $-$70.50272 & 0.65446 & 1.29496 & $-$53.96806 &$-$376.09815 &$-$219.19558 & 20.86879 & $-$168.44770 & $-$3339.75172 & $-$122.90226 & $-$129.01724 \\ OU & 553.75402 & 0.96886 & 38.30393 & 316.57714 & 736.48475 & 438.48980 & 59.35474 & 1061.50568 & 11142.81351 & 288.35969 & 90.00366 \\ \hline \multicolumn{12}{c}{Statistics of Sedna, 148209, 2004~VN$_{112}$, 2007~TG$_{422}$, 2010~GB$_{174}$, 2012~VP$_{113}$ and 2013~RF$_{98}$} \\ \hline Mean & 355.55661 & 0.83956 & 22.00602 & 105.81058 & 313.97529 & 59.78588 & 52.63982 & 658.47340 & 6904.16927 & 105.81058 & $-$46.02471 \\ Std. dev. & 110.14455 & 0.07477 & 5.60287 & 31.46021 & 20.47089 & 38.76357 & 18.27509 & 220.42583 & 3199.11104 & 31.46021 & 20.47089 \\ Median & 327.43521 & 0.85548 & 22.75598 & 112.91071 & 316.37528 & 38.59583 & 47.32201 & 607.54841 & 5921.13675 & 112.91071 & $-$43.62472 \\ Q$_{1}$ & 290.11545 & 0.82534 & 20.07888 & 79.16886 & 302.41767 & 28.68668 & 40.20255 & 521.87211 & 4954.24111 & 79.16886 & $-$57.58233 \\ Q$_{3}$ & 426.58491 & 0.87363 & 24.79958 & 129.50017 & 321.84310 & 90.33166 & 62.37693 & 794.83888 & 8908.43287 & 129.50017 & $-$38.15690 \\ IQR & 136.46947 & 0.04829 & 4.72070 & 50.33131 & 19.42543 & 61.64498 & 22.17438 & 272.96677 & 3954.19176 & 50.33131 & 19.42543 \\ OL & 85.41124 & 0.75290 & 12.99782 & 3.67190 & 273.27952 & $-$63.78080 & 6.94097 & 112.42195 & $-$977.04653 & 3.67190 & $-$86.72048 \\ OU & 631.28912 & 0.94607 & 31.88063 & 204.99713 & 350.98125 & 182.79914 & 95.63851 & 1204.28903 & 14839.72051 & 204.99713 & $-$9.01875 \\ \hline \end{tabular} \label{bary} \end{table*}
In this paper, we have re-analysed the various clusterings in ETNO orbital parameter space claimed in the literature and explored the visibility of Planet Nine within the context of improved constraints. Our results confirm the findings in Batygin \& Brown (2016) and Brown \& Batygin (2016) regarding clustering but using barycentric orbits. However, the observed overall level of clustering may not be maintained by a putative Planet Nine alone, other perturbers should exist. Summarizing: \begin{itemize} \item We confirm the existence of apsidal alignment and similar projected pole orientations among the currently known ETNOs. These patterns are consistent with the presence of perturbers beyond Pluto and/or, less likely, break-up of large asteroids at perihelion. \item If Planet Nine is at aphelion, it is most likely moving within $\alpha\in(3.0, 5.5)^{\rm h}$ and $\delta\in(-1, 6)\degr$ if $\Delta\varpi\sim$180\degr and $\Delta\Omega\sim$0\degr. \end{itemize}
16
7
1607.05633
The distribution of the orbital elements of the known extreme trans-Neptunian objects or ETNOs has been found to be statistically incompatible with that of an unperturbed asteroid population following heliocentric or, better, barycentric orbits. Such trends, if confirmed by future discoveries of ETNOs, strongly suggest that one or more massive perturbers could be located well beyond Pluto. Within the trans-Plutonian planets paradigm, the Planet Nine hypothesis has received much attention as a robust scenario to explain the observed clustering in physical space of the perihelia of seven ETNOs which also exhibit clustering in orbital pole position. Here, we revisit the subject of clustering in perihelia and poles of the known ETNOs using barycentric orbits, and study the visibility of the latest incarnation of the orbit of Planet Nine applying Monte Carlo techniques and focusing on the effects of the apsidal anti-alignment constraint. We provide visibility maps indicating the most likely location of this putative planet if it is near aphelion. We also show that the available data suggest that at least two massive perturbers are present beyond Pluto.
false
[ "barycentric orbits", "Planet Nine", "Monte Carlo techniques", "ETNOs", "trans-Neptunian", "trans-Plutonian", "Monte Carlo", "Pluto", "physical space", "orbital pole position", "much attention", "the apsidal anti-alignment constraint", "visibility maps", "perihelia", "the known extreme trans-Neptunian objects", "the trans-Plutonian planets paradigm", "poles", "aphelion", "the known ETNOs", "future discoveries" ]
7.592918
16.111055
-1
5012167
[ "Fu, Hai", "Hennawi, J. F.", "Prochaska, J. X.", "Mutel, R.", "Casey, C.", "Cooray, A.", "Kereš, D.", "Zhang, Z. -Y.", "Clements, D.", "Isbell, J.", "Lang, C.", "McGinnis, D.", "Michałowski, M. J.", "Mooley, K.", "Perley, D.", "Stockton, A.", "Thompson, D." ]
2016ApJ...832...52F
[ "The Circumgalactic Medium of Submillimeter Galaxies. I. First Results from a Radio-identified Sample" ]
13
[ "Department of Physics &amp; Astronomy, University of Iowa, Iowa City, IA 52242, USA", "Max-Planck-Institut fur Astronomie, Heidelberg, Germany", "Department of Astronomy and Astrophysics, UCO/Lick Observatory, University of California, 1156 High Street, Santa Cruz, CA 95064, USA", "Department of Physics &amp; Astronomy, University of Iowa, Iowa City, IA 52242, USA", "Department of Astronomy, the University of Texas at Austin, 2515 Speedway Blvd, Stop C1400, Austin, TX 78712, USA", "Department of Physics and Astronomy, University of California, Irvine, CA 92697, USA", "Department of Physics, Center for Astrophysics and Space Sciences, University of California at San Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ, UK ; ESO, Karl Schwarzschild Strasse 2, D-85748 Garching, Munich, Germany", "Astrophysics Group, Imperial College London, Blackett Laboratory, Prince Consort Road, London SW7 2AZ, UK", "Department of Physics &amp; Astronomy, University of Iowa, Iowa City, IA 52242, USA", "Department of Physics &amp; Astronomy, University of Iowa, Iowa City, IA 52242, USA", "Department of Physics &amp; Astronomy, University of Iowa, Iowa City, IA 52242, USA", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ, UK", "Oxford Centre For Astrophysical Surveys, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, USA", "Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 København Ø, Denmark", "Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu, HI 96822, USA", "Large Binocular Telescope Observatory, University of Arizona, 933 N. Cherry Ave, Tucson, AZ 85721, USA" ]
[ "2016MNRAS.463.3921M", "2017A&A...606A..17M", "2017ASSL..430..271F", "2017ApJ...844..123F", "2018ApJ...854..151M", "2018ApJ...859..125S", "2018ApJ...861..148M", "2018AsBio..18..630M", "2018AsBio..18..709C", "2020ApJ...905...86S", "2021ApJ...908..188F", "2023ApJ...957...91S", "2024arXiv240500795H" ]
[ "astronomy" ]
4
[ "galaxies: halos", "intergalactic medium", "quasars: absorption lines", "Astrophysics - Astrophysics of Galaxies" ]
[ "1985ApJ...298L...7H", "1986PASP...98..609H", "1992ARA&A..30..575C", "1995AJ....110.1526H", "1995PASP..107..375O", "1996ApJ...462..563N", "1997ApJ...490L...5S", "1998A&A...331L...1S", "1998Natur.394..241H", "1998Natur.394..248B", "1999ApJ...515..518E", "2001MNRAS.321..559B", "2002AJ....124.2364I", "2002Sci...295...82K", "2003AJ....125..465H", "2003PASP..115..389V", "2003PASP..115..688K", "2003SPIE.4841..962S", "2004ApJ...611..725B", "2004ApJ...617...64S", "2004MNRAS.349.1397C", "2004PASP..116..362C", "2005ApJ...622..772C", "2005MNRAS.359.1165G", "2005MNRAS.363....2K", "2005Natur.433..604D", "2006AJ....131....1H", "2006AJ....132..231P", "2006ApJ...650..592K", "2006ApJ...651...61H", "2006MNRAS.368....2D", "2006MNRAS.370.1057S", "2006MNRAS.372.1621C", "2006SPIE.6269E..4CE", "2007ASPC..376..127M", "2007ApJ...655..735H", "2007MNRAS.380..199I", "2008ApJ...680..246T", "2008MNRAS.383..615N", "2009ApJ...698.1010B", "2009ApJ...706L.173B", "2009ApJ...707.1201W", "2010A&A...518L...1P", "2010A&A...518L...3G", "2010A&A...518L...5N", "2010A&A...518L..23E", "2010A&A...518L..31I", "2010ApJ...712..942M", "2010ApJ...717..379B", "2010ApJ...718.1001B", "2010ApJ...725.1877B", "2010MNRAS.409...48R", "2010MNRAS.409...75H", "2010Natur.464..733S", "2011ApJ...737...67M", "2011ApJ...740...96H", "2011MNRAS.411.2739C", "2011MNRAS.412.1913I", "2011MNRAS.415.1479W", "2011MNRAS.415.2336R", "2011MNRAS.416..857S", "2011PASA...28..128C", "2012A&A...539A.155M", "2012A&A...541A..85M", "2012A&A...548A...4S", "2012ARA&A..50..455F", "2012ApJ...750...67R", "2012ApJ...752..152H", "2012ApJ...753..134F", "2012ApJ...761..140C", "2012ApJS..200....8K", "2012MNRAS.420..957Y", "2012MNRAS.421..284H", "2012MNRAS.424..933W", "2012MNRAS.424.1614O", "2013A&A...556A..55I", "2013ApJ...762L..19P", "2013ApJ...766...58H", "2013ApJ...767...88W", "2013ApJ...776..136P", "2013ApJS..204...19P", "2013MNRAS.429.3047B", "2013MNRAS.432.2012T", "2013Natur.498..338F", "2014ARA&A..52..589H", "2014ApJ...780...74F", "2014ApJ...782...68T", "2014ApJS..210...22V", "2014MNRAS.444.2870W", "2014PhR...541...45C", "2015ApJS..219...12A", "2015ApJS..221....2P", "2015MNRAS.449..987F", "2015MNRAS.452.2034R", "2015Natur.525..496N", "2016ApJ...823...17N", "2016MNRAS.461L..32F", "2016MNRAS.462.1989A", "2016MNRAS.462.3146V", "2017MNRAS.469.3396C" ]
[ "10.3847/0004-637X/832/1/52", "10.48550/arXiv.1607.00016" ]
1607
1607.00016_arXiv.txt
\label{sec:intro} \begin{deluxetable*}{rrrccc crrcc} \tablewidth{0pt} \tablecaption{VLA-observed \hers\ Sources \label{tab:vlaphoto}} \tablehead{ \colhead{Pair Name} & \colhead{RA$_{250}$} & \colhead{Dec$_{250}$} & \colhead{$S_{250}$} & \colhead{$S_{350}$} & \colhead{$S_{500}$} & \colhead{Int Time} & \colhead{RA$_{\rm 6 GHz}$} & \colhead{Dec$_{\rm 6 GHz}$} & \colhead{$S^{\rm peak}_{\rm 6 GHz}$} & \colhead{$S^{\rm int}_{\rm 6 GHz}$} \\ \colhead{} & \colhead{(deg)} & \colhead{(deg)} & \colhead{(mJy)} & \colhead{(mJy)} & \colhead{(mJy)} & \colhead{(min)} & \colhead{(deg)} & \colhead{(deg)} & \colhead{(\uJy/bm)} & \colhead{(\uJy)} \\ \colhead{(1)} & \colhead{(2)} & \colhead{(3)} & \colhead{(4)} & \colhead{(5)} & \colhead{(6)} & \colhead{(7)} & \colhead{(8)} & \colhead{(9)} & \colhead{(10)} & \colhead{(11)} } \startdata HeLMS 0015$+$0404 & 3.9286&$+$4.0715 & 61.7$\pm$6.0 & 67.8$\pm$5.7 & 54.3$\pm$7.2&26.5& 3.93038&$+$4.07262 & 69.4$\pm$6.7 & 72.4$\pm$13.4\\ HeLMS 0041$-$0410 & 10.3541&$-$4.1679 & 80.8$\pm$6.2 & 83.0$\pm$6.5 & 43.6$\pm$7.0&15.8& \nodata & \nodata & $<$35.1& \nodata \\ L6-XMM 0223$-$0605 & 35.8056&$-$6.0860 & 33.8$\pm$2.3 & 40.3$\pm$2.5 & 24.2$\pm$3.5&69.7& \nodata & \nodata & $<$14.4& \nodata \\ G09 0918$-$0039 &139.6159&$-$0.6644 & 41.1$\pm$6.9 & 49.7$\pm$8.1 & 31.2$\pm$9.1&18.9& \nodata & \nodata & $<$17.7& \nodata \\ G09 0920$+$0024 &140.2475&$+$0.4049 & 35.0$\pm$7.0 & 51.6$\pm$8.1 & 32.0$\pm$8.9&22.9& \nodata & \nodata & $<$18.6& \nodata \\ NGP 1313$+$2924 &198.4530&$+$29.4126& 59.7$\pm$5.6 & 78.5$\pm$6.6 & 53.6$\pm$7.8&17.2& \nodata & \nodata & $<$42.3& \nodata \\ NGP 1330$+$2540 &202.5866&$+$25.6749& 49.2$\pm$5.8 & 54.3$\pm$6.4 & 29.3$\pm$7.8&18.0& \nodata & \nodata & $<$15.3& \nodata \\ NGP 1333$+$2357 &203.3743&$+$23.9592& 30.4$\pm$5.4 & 31.8$\pm$6.4 & 29.1$\pm$7.5&55.2&203.37514&$+$23.95909 & 20.0$\pm$3.0 & 31.0$\pm$8.1 \\ NGP 1335$+$2805 &203.9409&$+$28.0986& 41.7$\pm$5.5 & 49.8$\pm$6.4 & 38.2$\pm$7.7&27.6&203.94249&$+$28.09750 & 38.3$\pm$4.5 & 51.4$\pm$11.0\\ G15 1413$+$0058 &213.4580&$+$0.9725 & 46.6$\pm$6.4 & 52.8$\pm$7.7 & 36.1$\pm$8.5&16.2&213.45743&$+$0.97321 & 37.3$\pm$5.8 & 37.3$\pm$11.6\\ G15 1435$+$0110 &218.9043&$+$1.1682 & 63.0$\pm$6.7 & 63.8$\pm$8.0 & 56.6$\pm$8.8& 9.2&218.90494&$+$1.16958 & 75.6$\pm$7.3 & 89.1$\pm$16.0\\ G15 1450$+$0026 &222.6773&$+$0.4351 & 47.5$\pm$6.9 & 47.9$\pm$8.1 & 29.1$\pm$8.9&16.3& \nodata & \nodata & $<$18.0& \nodata \\ L6-FLS 1712$+$6001 &258.0352&$+$60.0281& 32.3$\pm$2.2 & 34.0$\pm$2.4 & 22.4$\pm$3.6&30.1&258.03111&$+$60.02722 & 10.8$\pm$2.1 & 10.8$\pm$4.2 \\ & & & & & & & 258.04010 & $+$60.02625 & 13.9$\pm$2.1 & 15.6$\pm$4.5 \enddata \tablecomments{ L6-FLS 1712$+$6001 has two radio counterparts (see Fig.~\ref{fig:detections}); the first line shows the nominal counterpart, which is closer to the \hers\ position although slightly fainter. Columns (2-6) list the \hers\ 250~\um\ positions and the photometry at 250, 350, and 500~\um. Column (7) is the total VLA on-source integration time. Columns (8-9) list the positions of the radio counterparts. Columns (10-11) are the peak flux density in \uJy/bm and the integrated flux density in \uJy, both of which are derived by fitting an elliptical Gaussian model to the source. The uncertainty of peak flux density is given by the rms noise in the map at the source position, while the uncertainty of the integrated flux density is estimated using the formulae provided by \citet{Hopkins03a}, which includes the 1\% uncertainty in the VLA flux-density scale at 6~GHz \citep{Perley13}. } \end{deluxetable*} The first milli-Jansky-level submillimeter surveys discovered a population of distant submillimeter-bright galaxies (SMGs), namely, unresolved sources with 850~\um\ flux density ($S_{850}$) greater than 3$-$5~mJy \citep{Smail97,Barger98,Hughes98,Eales99}. The SMGs selected at wavelengths between 850~\um\ and 1~mm are intense starbursts (SFR $\gtrsim 500$~\msunyr) at a median redshift of $z \sim 2.5$ \citep{Chapman05,Wardlow11,Yun12,Smolcic12}. The intense star formation is dust-enshrouded so that the SMGs radiate most of their bolometric luminosity in the far-infrared (IR). The observed molecular and stellar emission indicates that they are massive gas-rich galaxies \citep[\hl{\mbox{$M_{\rm gas} \sim M_{\rm star} \sim 10^{11}$\,\msun}}; e.g., ][]{Michalowski10a,Hainline11,Bothwell13}, but the typical halo mass of SMGs remains uncertain, with estimates ranging from $10^{12}$ to $10^{13}$\,\msun. Two lines of evidence suggest that SMGs may inhabit dark matter halos as massive as $\sim10^{13}$~\msun: (1) their strong clustering strength estimated from either the angular two-point correlation function \citep[e.g.,][]{Scott06,Weis09} or the cross-correlation function between SMGs and other high-redshift galaxies \citep[e.g.,][]{Hickox12}, and (2) their high stellar mass and the \ms$-$\mh\ relation from abundance matching \citep[\mh~$= 6\times10^{12}$~\msun\ for \ms~$= 10^{11}$~\msun\ at $z = 2$; e.g.,][]{Behroozi10}. However, because source blending due to the large beams of single-dish (sub)millimeter telescopes may have significantly elevated clustering strength \citep{Cowley16} and the stellar mass estimates remain uncertain within an order of magnitude \citep[e.g.,][]{Hainline11,Michalowski12, Targett12}, it is possible that a typical SMG may inhabit much less massive halos ($\sim$10$^{12}$\,\msun). SMGs are absent in the local universe and it is commonly thought that they have evolved into the massive ellipticals today \citep[e.g.,][]{Blain04a,Toft14}. To understand the evolution of SMGs, it is imperative to know how long the observed intense star formation would last. For $10^{13}$~\msun\ halos at $z = 2.5$, the average baryonic accretion rate from the mass growth rate of dark matter haloes is $\dot{M}_{\rm gas} \equiv 0.18 \times \dot{M}_{\rm halo} \simeq 1.4\times10^3$~\msunyr\ \citep{Neistein08,Bouche10}. In such massive halos, it is expected that most of the baryons will be shock-heated to the virial temperature of the halo ($\sim10^7$~K) so that only a small fraction of the accreted gas can actually cool and accrete onto galaxies \citep[e.g.,][]{Keres05, Dekel06}. Therefore, the ongoing gas accretion is unlikely to sustain the extreme SFRs. Without a comparable gas supply rate, the SFR would decline with an $e$-folding timescale of only $\sim$200~Myr \citep[2\mg/SFR; e.g.,][]{Greve05, Tacconi08, Ivison11, Bothwell13, Fu13}. At such a rate, the SMGs would become red sequence galaxies in only a Gyr or 5 $e$-folding times\footnote{This is the time it would take to decrease the specific SFR (SFR/\ms) from $\sim10^{-9}$~Gyr$^{-1}$ for the SMGs at the observed epoch to $\sim10^{-11}$~Gyr$^{-1}$ for the red sequence at $z \sim 2$ \citep{Brammer09}.}. Such a short transitional time of a significant high-redshift star-forming population might help explain the rapid build-up of the massive end of the red sequence at $z > 1$ \citep[e.g.,][]{Ilbert13}. Starbursts thus provide an alternative mechanism to the QSO-mode feedback \citep[e.g.][]{Silk98,Di-Matteo05} to form red and dead galaxies. Note that both mechanisms still require feedback from radio jets \citep[i.e., the maintenance mode; e.g.,][]{Fabian12,Heckman14} to prevent the hot gaseous halo from cooling. On the other hand, if SMGs were in $10^{12}$~\msun\ halos, the intense star formation is also unsustainable because the gas accretion rate is only $\sim$110~\msunyr\ at $z = 2.5$. However, could there be enough cool gas in the circumgalactic medium (CGM) around SMGs to fuel a prolonged starburst phase \citep[e.g. see the simulation of][]{Narayanan15}? The CGM of co-eval QSOs may give us a hint, because they inhabit comparably massive ($\sim10^{12.6}\,M_{\odot}$) halos \citep{White12}. Contrary to the expected dominance of virialized X-ray plasma, absorption line spectroscopy of a statistical sample of $z \sim 2$ projected QSO pairs reveals the prevalence of cool ($T \sim 10^4$~K), metal-enriched ($Z \geq 0.1 Z_\odot$), and optically thick \lya\ absorbers ($N_{\rm HI} \geq 10^{17.2}$~\cmsq) extending to at least the expected virial radius of 160~kpc (the ``QSO Probing QSO'' [QPQ] project: \citealt{Hennawi06a,Hennawi07,Prochaska13,Prochaska13a}). The high observed covering factor of the cool CGM gas ($\gtrsim 60\%$) in $\sim10^{12.6}\,M_{\odot}$ halos has been compared to predictions from numerical simulations. While several studies found that they cannot reproduce the high covering factor around QSOs (\citealt{Fumagalli14, Faucher-Giguere15}, but see \citealt{Rahmati15}), it has been argued that efficient star-formation-driven winds from accreted satellite galaxies that interact with cosmological filaments are required to increase the \HI\ covering factor to the observed level, which is only resolved in the highest resolution cosmological zoom simulations \citep{Faucher-Giguere16}. These high-resolution simulations predict that the covering factor is roughly independent of SFR, decreases with redshift, and has relatively large halo-halo variations. In the relevant halo mass range, $\sim 3\times 10^{12}-10^{13}$\,\msun, simulations show \HI\ covering factors of $\sim 30-80\%$ with little mass dependence at $z \sim 2$ \hl{\mbox{\citep{Faucher-Giguere16}}}. Given the estimated halo masses for QSOs and SMGs, one could expect similar covering factors if these are determined primarily by the interplay between gas infall and star-formation-driven outflows. To test this, we exploit QSO absorption line spectroscopy to probe the CGM of SMGs. We first present the data sets and the method we used to select projected \sqps\ in \S~\ref{sec:sample}. We then describe our followup observations in \S~\ref{sec:obs}, including radio interferometer imaging, near-infrared spectroscopy, and optical spectroscopy. We present our analysis and results in \S~\ref{sec:result}, including a comparison between the covering factor of optically thick gas around SMGs and that of $z \sim 2$ QSOs. We summarize the results and conclude in \S~\ref{sec:summary}. Throughout we adopt a $\Lambda$CDM cosmology with $\Omega_{\rm m}=0.27$, $\Omega_\Lambda=0.73$ and $H_0$ = 70 km~s$^{-1}$~Mpc$^{-1}$.
\label{sec:summary} Motivated by the unique properties of SMGs and their purported evolutionary link to high-redshift QSOs and today's massive ellipticals, we have started a project to use QSO absorption line spectroscopy to probe the diffuse cool \HI\ gas in the CGM of SMGs. This work requires a sample of projected \sqps, which are extremely rare. Wide-area sub-millimeter surveys are needed to compile such a sample. Thanks to the advent of \hers, we have identified 163 \sqps\ with bright $z > 2.5$ QSOs ($g_{\rm QSO} < 22$) and angular separations between 5\arcsec\ and 36\arcsec\ from a suite of wide-area \hers\ surveys and spectroscopic QSO surveys. Extensive followup observations are required to carry out the absorption line study. To allow slit spectroscopy, the first stage is to use an interferometer to pin down the positions of the \hers\ sources. This paper focuses on a subsample of 13 \sqps\ that were observed with the VLA in C-band. With an average integration time of 25~min per source, the VLA detected sources within the \hers\ beam in six fields. One of the six \sqp\ turns out to be a far-IR-luminous QSO at $z = 2.973$, while the \hers\ source in another pair turns out to be a pair of SMGs. Hence, we effectively identified six \sqps\ from observations of 13 fields (46\%). The second stage is to measure the spectroscopic redshifts of the SMGs. We observed five of the six VLA-identified \sqps\ with near-IR spectroscopy. We were able to determine the redshifts for three of the five SMGs (60\%) from the redshifted H$\alpha$ and \NII\ lines. The remaining two sources show only featureless continuum in the near-IR windows, making it impossible to determine an accurate redshift. The last stage is to obtain optical spectroscopy of the QSOs once it is confirmed that the QSOs are in the background of the SMGs. \hl{Because we have selected only the brightest QSOs, the success rate of this step is approaching 100\%}. Because of the low success rates of the first two followup stages, only $\sim$23\% of the initial sample (i.e., three \sqps\ out of 13) are spectroscopically confirmed and are suitable for final absorption line study. Our main findings are as follows: \begin{enumerate} \item The near-IR spectra of the five VLA-detected sources are all detected in hour-long integrations with 8-meter telescopes, although only three of those show emission lines that yielded accurate spectroscopic redshifts. Since we positioned the slits on the VLA positions, this result shows that the spatial offset, if any, between the near-IR and the radio counterparts is less than an arcsec (the slit width), consistent with the finding from differential lensing in strongly lensed sources \citep[e.g.,][]{Fu12b}. \item The VLA-identified \hers\ ``350~\um peakers'' at $2.0 < z < 2.6$ are similar to SMGs selected at longer wavelengths (i.e., 850~\um\ to 1~mm), in terms of the \NII/\Ha\ ratio (i.e., gas metallicity), the IR luminosity (i.e., SFR), and the IR-to-radio luminosity ratio (i.e., radio excess due to AGN). The VLA-identified SMGs are optically faint and unbiased to radio-loud AGNs, so they indeed represent a galaxy population distinct from the optical selected QSOs and the Lyman break galaxies (LBGs) in the same redshift range. \item Strong \HI\ \lya\ absorption is found in the background QSO spectra for all of the three spectroscopically confirmed \sqps\ with impact parameters between $100 < R_\bot < 200$~kpc. Here we have adopted a much narrower search window ($\pm$600~\kms) than the QPQ study, further reducing the level of contamination from physically unrelated clouds. However, none of the three absorption line systems seems optically thick at the Lyman limit (i.e., LLSs with $N_{\rm HI} > 10^{17.2}$~cm$^{-2}$), in contrast to the $\sim$60\% covering factor of LLSs around QSOs from the QPQ study despite similar data quality, foreground redshifts, and impact parameters. \end{enumerate} Our comparison thus suggests {\it either} that SMGs do not have a substantial neutral gas reservoir in their halos that could potentially fuel a prolonged star formation phase {\it or} that SMGs inhabit $\sim10^{12}$\,\msun\ halos so that our sightlines are yet to probe inside their virial radii. If the latter, their halos are comparable to those of LBGs. \citet{Rudie12} found an optically thick covering factor of $30\pm14$\% around LBGs at $z \sim 2.3$ and $R_\bot < 90$~kpc. Note that this is $\sim$2$\times$ lower than that of coeval QSOs and is consistent with the 1$\sigma$ confidence interval that we were able to place for the SMGs. On the other hand, the difference in the optically thick \HI\ covering factor between SMGs and QSOs casts doubt on the evolutionary link between the two populations, unless AGN outflows can somehow affect the physical state of gas at hundreds of kpc scales within its short lifetime. Our final conclusion is limited by the small sample size. To enable a more robust comparison with previous absorption-line studies, we badly need to increase the effective yield of our survey from the current level of $\sim$23\%. In a future publication, we will present observations with the Atacama Large Millimeter/submillimeter Array (ALMA) to pinpoint the positions of the \hers\ sources in the \sqps. Observing at a wavelength (870~\um) much closer to the selection wavelengths (250 to 500~\um), we expect doubling the detection rate with integration times of just several minutes per source. The \hers\ sources in our sample are too faint to allow a quick CO redshift search with ALMA \citep[e.g., the survey of strongly lensed SMGs by][]{Weis13}, but spectrographs on large optical telescopes covering the full optical$+$IR range at a moderate spectral resolution ($R \gtrsim 2000$) will greatly increase the redshift search range and decrease the areas blinded by strong airglow lines.
16
7
1607.00016
We present the first results from an ongoing survey to characterize the circumgalactic medium (CGM) of massive high-redshift galaxies detected as submillimeter galaxies (SMGs). We constructed a parent sample of 163 SMG-QSO pairs with separations less than ∼36″ by cross-matching far-infrared-selected galaxies from Herschel with spectroscopically confirmed QSOs. The Herschel sources were selected to match the properties of the SMGs. We determined the sub-arcsecond positions of six Herschel sources with the Very Large Array and obtained secure redshift identification for three of those with near-infrared spectroscopy. The QSO sightlines probe transverse proper distances of 112, 157, and 198 kpc at foreground redshifts of 2.043, 2.515, and 2.184, respectively, which are comparable to the virial radius of the ∼10<SUP>13</SUP> M <SUB>⊙</SUB> halos expected to host SMGs. High-quality absorption-line spectroscopy of the QSOs reveals systematically strong H I Lyα absorption around all three SMGs, with rest-frame equivalent widths of ∼2-3 Å. However, none of the three absorbers exhibit compelling evidence for optically thick H I gas or metal absorption, in contrast to the dominance of strong neutral absorbers in the CGM of luminous z ∼ 2 QSOs. The low covering factor of optically thick H I gas around SMGs tentatively indicates that SMGs may not have as prominent cool gas reservoirs in their halos as the coeval QSOs and that they may inhabit less massive halos than previously thought.
false
[ "SMGs", "prominent cool gas reservoirs", "strong neutral absorbers", "submillimeter galaxies", "secure redshift identification", "systematically strong H I Lyα absorption", "foreground redshifts", "compelling evidence", "massive high-redshift galaxies", "optically thick H I gas or metal absorption", "rest-frame equivalent widths", "CGM", "∼2", "less massive halos", "none", "Herschel", "the coeval QSOs", "the SMGs", "contrast", "High-quality absorption-line spectroscopy" ]
13.447112
7.376157
132
12472788
[ "Whittam, I. H.", "Riley, J. M.", "Green, D. A.", "Jarvis, M. J." ]
2016MNRAS.462.2122W
[ "The faint source population at 15.7 GHz - III. A high-frequency study of HERGs and LERGs" ]
21
[ "Physics and Astronomy Department, University of the Western Cape, Bellville 7535, South Africa", "Astrophysics Group, Cavendish Laboratory, 19 JJ Thomson Avenue, Cambridge CB3 0HE, UK", "Astrophysics Group, Cavendish Laboratory, 19 JJ Thomson Avenue, Cambridge CB3 0HE, UK", "Physics and Astronomy Department, University of the Western Cape, Bellville 7535, South Africa; Astrophysics, University of Oxford, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, UK" ]
[ "2017MNRAS.464.3357W", "2017MNRAS.471..908W", "2018A&A...609A...1B", "2018AstBu..73..142Z", "2018MNRAS.474.4133H", "2018MNRAS.480..358W", "2019A&A...621A..19R", "2019Galax...7...76B", "2019MNRAS.482.2294B", "2020MNRAS.493.2841W", "2020MNRAS.494.2053P", "2021A&A...650A.127T", "2021A&ARv..29....3O", "2021MNRAS.503.3492T", "2022A&A...662A..28M", "2022JApA...43...97S", "2022MNRAS.509.2150H", "2022MNRAS.515.5539T", "2022MNRAS.516..245W", "2023A&ARv..31....3B", "2024JHEAp..42...21G" ]
[ "astronomy" ]
4
[ "surveys", "galaxies: active", "radio continuum: galaxies", "Astrophysics - Astrophysics of Galaxies" ]
[ "1973A&A....24..337S", "1974MNRAS.167...31S", "1974MNRAS.167P..31F", "1979MNRAS.188..111H", "1990AJ.....99...14B", "1993ARA&A..31..473A", "1994ASPC...54..201L", "1994ApJS...95....1E", "1995ApJ...452..710N", "1995PASP..107..803U", "1997ApJ...475..479W", "1997MNRAS.286..241J", "1998AJ....115.1693C", "1998PASP..110..493O", "2000A&A...363..887D", "2001MNRAS.326.1279H", "2003MNRAS.346.1055K", "2003PASP..115..897L", "2004ApJ...602..116W", "2004ApJS..154..166L", "2004MNRAS.352..909C", "2005ApJ...621..256S", "2005MNRAS.362....9B", "2005MNRAS.362...25B", "2005Natur.436..666M", "2006ApJ...642...96E", "2006ApJ...647..161O", "2006MNRAS.367..323B", "2006MNRAS.370.1556B", "2006MNRAS.370.1893H", "2006MNRAS.371..963B", "2007AJ....133..186L", "2007MNRAS.376.1849H", "2007MNRAS.379..894B", "2007MNRAS.379.1599L", "2007MNRAS.381..211S", "2007MNRAS.381..589M", "2008A&A...479..283B", "2008A&A...490..893T", "2008AJ....136.1889O", "2008MNRAS.387.1037G", "2008MNRAS.391.1545Z", "2009ASPC..402..221S", "2009ApJS..185..433W", "2010A&A...518L..20V", "2010BASI...38..103G", "2010MNRAS.402.2403M", "2010MNRAS.406.1841H", "2010MNRAS.408.1187H", "2011ApJ...740...20P", "2011MNRAS.410.1360H", "2011MNRAS.415.2699A", "2011MNRAS.415.2708A", "2011MNRAS.416..927G", "2011MNRAS.417.2651M", "2012ApJ...748..142D", "2012ApJS..198....1F", "2012MNRAS.421.1569B", "2012PASP..124..714M", "2013A&A...551A..97M", "2013MNRAS.428.1958R", "2013MNRAS.429.2080W", "2013MNRAS.429.2407H", "2013MNRAS.432..609V", "2013MNRAS.436.3759B", "2014MNRAS.438..796S", "2014MNRAS.438.1149G", "2014MNRAS.440...40W", "2014MNRAS.445..955B", "2015A&A...576A..38B", "2015ApJS..219...12A", "2015MNRAS.447.1184F", "2015MNRAS.450.1538W", "2015MNRAS.453.4244W", "2015fers.confE..27V", "2016MNRAS.457..730P" ]
[ "10.1093/mnras/stw1725", "10.48550/arXiv.1607.03709" ]
1607
1607.03709_arXiv.txt
\label{section:intro} In the first two papers in this series (\citealt{2013MNRAS.429.2080W,Paper II}, referred to as Papers I and II respectively) we studied the properties of a sample of sources selected from the Tenth Cambridge (10C; \citealt{2011MNRAS.415.2708D,2011MNRAS.415.2699F}) survey at 15.7 GHz in the Lockman Hole. The 10C survey is complete to 0.5~mJy, making it the deepest high-frequency radio survey to date. In Paper II we found that the vast majority ($\geqslant 94$ per cent) of the sources in the 10C sample are radio galaxies; the 10C sample is therefore the ideal starting point for a study of the properties of faint radio galaxies. It is well known that powerful extended radio galaxies can be split into two classes according to their morphology following the Fanaroff--Riley (FR) scheme; FRI sources have the highest surface brightness near the core, while FRII sources are brightest near the lobe edges and have more collimated jets \citep{1974MNRAS.167P..31F}. There is also a divide in luminosity between the two classes, with the more powerful FRII sources having $L_{1.4~\rm GHz} > 10^{24.5}~\rm W \, Hz^{-1}$, while the FRI sources predominantly have luminosities below this value. Some recent papers (e.g.\ \citealt{2014MNRAS.438..796S}, \citealt{2015A&A...576A..38B}) have introduced the `FR0' classification to describe radio galaxies which lack the extended emission typical of FRI and FRII sources. There are two main types of compact radio galaxies; compact-steep spectrum sources (CSS), which have linear sizes between 1 and 15~kpc and convex spectra which peak below $\sim 500$~MHz, and gigahertz-peaked spectrum sources (GPS), which are smaller with linear sizes $\lesssim 1$~kpc and have spectra which peak between 500~MHz and 10~GHz (see \citealt{1998PASP..110..493O} for a review). Both classes of compact sources are powerful, with typical luminosities $L_{1.4~\rm GHz} \gtrsim 10^{25}~\rm W \, Hz^{-1}$. There is evidence that both GPS and CSS sources are young radio galaxies (\citealt{2008arXiv0802.1976S} and references therein). It has been known for some time \citep{1979MNRAS.188..111H} that the properties of radio galaxies are not fully explained by the conventional model of an AGN, consisting of an accretion disk surrounded by a dusty torus \citep{1993ARA&A..31..473A}. Based on this model, we would expect radio-loud objects viewed close to the jet axis to show both broad and narrow optical emission lines, while radio-loud objects viewed at larger angles to the jet would only show narrow lines and would have a clear mid-infrared signature of the dusty torus. However, many radio-loud AGN lack the expected narrow-line optical emission and do not display evidence of an obscuring torus (e.g.\ \citealt{2004ApJ...602..116W}). Subsequent studies have suggested that there are two fundamentally distinct accretion modes, known as `cold mode' and `hot mode' (see \citealt{2005MNRAS.362...25B,2007MNRAS.376.1849H}) which could be responsible for these differences (these modes are sometimes referred to as `quasar' and `radio' modes respectively). Cold-mode accretion occurs when cold gas is accreted onto the central black hole through a radiatively efficient, geometrically thin, optically thick accretion disk (e.g.\ \citealt{1973A&A....24..337S}) and gives rise to the traditional model of an AGN \citep{1993ARA&A..31..473A}. These objects therefore show high-excitation lines in their optical spectra, so are often referred to as high-excitation radio galaxies (HERGs). `Hot mode' sources, however, are thought to be fuelled by the accretion of warm gas through advection-dominated accretion flows (e.g.\ \citealt{1995ApJ...452..710N}) and lack many of the typical signatures of AGN, such as strong optical emission lines. These objects are often referred to as low-excitation radio galaxies (LERGs). LERGs typically show no evidence for a dusty torus \citep{2006ApJ...647..161O} or for accretion-related X-ray emission \citep{2006MNRAS.370.1893H}. There are also differences in the host galaxies of the two populations, with HERGs found in less massive and bluer galaxies than LERGs \citep{2005MNRAS.362....9B,2008A&A...490..893T,2010MNRAS.406.1841H,2011MNRAS.410.1360H,2015MNRAS.450.1538W}. Although both HERGs and LERGs are found across the full range in radio luminosities, LERGs seem to dominate at lower luminosities and HERGs at higher luminosities \citep{2012MNRAS.421.1569B}. There is a substantial overlap between the FRI/II classification and the HERG and LERG classes, with most FRI sources being LERGs and most FRIIs being HERGs \citep{2012MNRAS.421.1569B}. There are, however, differences in the two classifications, in particular a significant population of FRII LERGs has been found \citep{1994ASPC...54..201L}. \citet{2012MNRAS.421.1569B} showed that low-redshift HERGs and LERGs have distinct accretion rates; HERGs typically accrete at between 1 and 10 per cent of their Eddington rate, while LERGs generally have accretion rates much less than 1 per cent of their Eddington rate. \citet{2015MNRAS.447.1184F} also find evidence for this at $z \sim 1$ using mid-infrared data. Therefore a picture is emerging where HERGs accrete cold gas at a relatively high rate \citep{2012MNRAS.421.1569B} and radiate efficiently across the whole electromagnetic spectrum (e.g.\ \citealt{1994ApJS...95....1E}). This causes them to produce a stable accretion disc and therefore display the typical properties of an AGN. The cold gas leads to star formation \citep{2013MNRAS.432..609V,2013MNRAS.429.2407H}, causing the host galaxies of HERGs to be relatively blue (e.g.\ \citealt{2003MNRAS.346.1055K}). In addition, HERGs are more prevalent at earlier cosmic epochs, where higher rates of mergers and interactions provided a steady supply of cold gas. LERGs, however, slowly accrete warm gas from the X-ray emitting halo of the galaxy or cluster. They radiate inefficiently, emitting the bulk of their energy in kinetic form as powerful jets (e.g.\ \citealt{2007MNRAS.381..589M}). They therefore tend to be hosted by massive galaxies, often at the centres of groups or clusters \citep{1990AJ.....99...14B,2007MNRAS.379..894B}. These galaxies have an old, passive stellar population and show little cosmic evolution out to $z \sim 1$ \citep{2004MNRAS.352..909C,2007MNRAS.381..211S,2014MNRAS.445..955B}. The fundamentally different accretion modes of HERGs and LERGs cause them to have different radio properties. \citet{2011MNRAS.417.2651M} studied the properties of high flux density ($S_{20~\rm GHz} > 40$~mJy) HERGs and LERGs in the Australia Telescope 20~GHz (AT20G; \citealt{2010MNRAS.402.2403M}) sample and found that while both accretion modes display a range of radio properties, a higher fraction of HERGs are extended and have steep spectra. They suggest that this is because HERGs are accreting more efficiently than LERGs, and therefore have a greater chance of producing more luminous jets and lobes. They also find that HERGs display different properties depending on their orientation, with objects displaying broad emission lines tending to be flat spectrum, and objects with narrow lines tending to be steep spectrum, as predicted by orientation models (\citealt{1993ARA&A..31..473A,1995PASP..107..803U}). LERGs, however, display no orientation effects. These results are consistent with the scenario in which HERGs have a typical AGN accretion disk and torus while LERGs do not. Most studies of HERGs and LERGs have focussed on high flux density (e.g.\ \citealt{2011MNRAS.417.2651M}) or low redshift (e.g.\ \citealt{2012MNRAS.421.1569B}) sources, with the exception of \citet{2015MNRAS.447.1184F}, who studied a small sample at $z \sim 1$. In this work, we focus on a sample of high-frequency selected mJy and sub-mJy sources. In Section~\ref{section:radio-data} we outline the data used in this study. In Section~\ref{section:classifications} we describe how this data is used to separate the radio galaxies into HERGs and LERGs and the properties of these HERGs and LERGs are explored in Section~\ref{section:object_properties}. In Section~\ref{section:FR-sources} the sample is split into different radio morphological classes and the properties of these classes are discussed and compared to the HERG and LERGs classifications. The sample is compared to the higher-flux density AT20G sample in Section~\ref{section:other-work} and the conclusions are presented in Section~\ref{section:chap5-conclusions}.
\label{section:chap5-conclusions} We have studied the properties of a sample of 96 radio galaxies selected from the 10C sample in the Lockman Hole. To distinguish between high and low-excitation radio galaxies (HERGs and LERGs) three different methods are used; optical compactness, X-ray observations and mid-infrared colour--colour diagrams. These methods are combined to produce overall HERG and LERG classifications; a total of 32 sources are classified as `probable HERGs', 35 as `probable LERGs' and 29 remain unclassified. 17 sources are also classified using their optical spectra; the spectroscopic classifications agree with the classifications derived here in 94 per cent of cases, showing that this is a reliable way of distinguishing between HERGs and LERGs. The properties of these HERGs and LERGS are then compared. We find that the HERGs in our sample tend to be found at higher redshifts, have flatter spectra, higher 15.7-GHz flux densities and smaller linear sizes than LERGs. Note however that the relatively large proportion of unclassified sources have the potential to change this result. This result is in contrast to that found for the higher-flux-density AT20G sample, where the HERGs have steeper spectra and are more extended than the LERGs. This suggests that the HERGs in the 10C sample have different properties to their higher-flux-density counterparts, lacking the powerful extended emission typical of FRI and FRII sources, and instead being dominated by their cores. Low-frequency (610~MHz) radio images, along with radio spectral indices and linear sizes, are used to split the sources into different radio morphological classes. Although 18 sources are found to be FRI or FRII sources and 13 others are significantly extended, the majority of the sample (65 sources) do not display any extended emission on arcsecond scales and are therefore classified as FR0 sources. These FR0 sources are further subdivided into candidate GPS sources (13 sources) and candidate CSS sources (10 sources), while the remaining 42 sources could not be classified. The sub-sample of fainter 10C sources with 15.7-GHz flux density $<$~1~mJy contains a higher proportion of compact FR0 sources (81 per cent), consistent with the result from Paper I that the majority of the faint 10C sources have flat radio spectra. There is some indication that the CSS and GPS sources are more likely to be HERGs than LERGs, supporting the idea that the 10C HERGs tend to be smaller and in some cases more core-dominated than the LERGs, but a larger proportion of the source must be classified optically into HERGs and LERGs before this conclusion can be confirmed. By comparing these results to those from both the full AT20G survey \citep{2011MNRAS.417.2651M} and the sub-sample of nearby ($z \sim 0.05$) sources in the AT20G-6dFGS sample \citep{2014MNRAS.438..796S} we have studied the high-frequency extragalactic source population over a wide range in flux density. We find that the fainter 10C sources are not simply higher-redshift versions of the brighter AT20G sources; the two samples are found at similar redshifts, but the 10C sources have significantly lower luminosities. The nature of the optical counterparts to the radio galaxies changes with flux density; at high flux densities most radio sources are associated with quasars, while at low flux densities the optical counterparts are primarily galaxies. We find evidence that the 10C sources may be higher-redshift versions of the lower-luminosity, compact radio galaxies found in the AT20G-6dFGS sub-sample. This work shows that the faint, flat spectrum radio galaxies which dominate the high-frequency radio sky below 1~mJy are a mixed population of HERGs and LERGs, most of which lack the extended emission typical of FRI or II radio galaxies and have lower luminosities than the higher-flux density sources found in shallower surveys.
16
7
1607.03709
A complete sample of 96 faint (S &gt; 0.5 mJy) radio galaxies is selected from the Tenth Cambridge (10C) survey at 15.7 GHz. Optical spectra are used to classify 17 of the sources as high-excitation or low-excitation radio galaxies (HERGs and LERGs, respectively), for the remaining sources three other methods are used; these are optical compactness, X-ray observations and mid-infrared colour-colour diagrams. 32 sources are HERGs and 35 are LERGs while the remaining 29 sources could not be classified. We find that the 10C HERGs tend to have higher 15.7-GHz flux densities, flatter spectra, smaller linear sizes and be found at higher redshifts than the LERGs. This suggests that the 10C HERGs are more core dominated than the LERGs. Lower-frequency radio images, linear sizes and spectral indices are used to classify the sources according to their radio morphology; 18 are Fanaroff and Riley type I or II sources, a further 13 show some extended emission, and the remaining 65 sources are compact and are referred to as FR0 sources. The FR0 sources are sub-divided into compact, steep-spectrum sources (13 sources) or gigahertz-peaked spectrum sources (10 sources) with the remaining 42 in an unclassified class. FR0 sources are more dominant in the subset of sources with 15.7-GHz flux densities &lt;1 mJy, consistent with the previous result that the fainter 10C sources have flatter radio spectra. The properties of the 10C sources are compared to the higher-flux density Australia Telescope 20 GHz (AT20G) survey. The 10C sources are found at similar redshifts to the AT20G sources but have lower luminosities. The nature of the high-frequency selected objects changes as flux density decreases; at high flux densities the objects are primarily quasars, while at low flux densities radio galaxies dominate.
false
[ "sources", "FR0 sources", "low flux densities", "flux density", "high flux", "radio galaxies", "flatter radio spectra", "higher redshifts", "gigahertz-peaked spectrum sources", "the AT20G sources", "the remaining sources", "smaller linear sizes", "linear sizes", "The FR0 sources", "lower luminosities", "higher 15.7-GHz flux densities", "10 sources", "10C sources", "13 sources", "32 sources" ]
14.894485
5.761029
119
2774734
[ "Zhao, Gong-Bo", "Wang, Yuting", "Saito, Shun", "Wang, Dandan", "Ross, Ashley J.", "Beutler, Florian", "Grieb, Jan Niklas", "Chuang, Chia-Hsun", "Kitaura, Francisco-Shu", "Rodriguez-Torres, Sergio", "Percival, Will J.", "Brownstein, Joel R.", "Cuesta, Antonio J.", "Eisenstein, Daniel J.", "Gil-Marín, Héctor", "Kneib, Jean-Paul", "Nichol, Robert C.", "Olmstead, Matthew D.", "Prada, Francisco", "Rossi, Graziano", "Salazar-Albornoz, Salvador", "Samushia, Lado", "Sánchez, Ariel G.", "Thomas, Daniel", "Tinker, Jeremy L.", "Tojeiro, Rita", "Weinberg, David H.", "Zhu, Fangzhou" ]
2017MNRAS.466..762Z
[ "The clustering of galaxies in the completed SDSS-III Baryon Oscillation Spectroscopic Survey: tomographic BAO analysis of DR12 combined sample in Fourier space" ]
57
[ "National Astronomy Observatories, Chinese Academy of Science, Beijing, 100012, P.R.China; Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "National Astronomy Observatories, Chinese Academy of Science, Beijing, 100012, P.R.China; Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Starße 1, D-85740 Garching bei München, Germany; Kavli Institute for the Physics and Mathematics of the Universe (IPMU), The University of Tokyo, Kashiwa, Chiba 277-8583, Japan", "National Astronomy Observatories, Chinese Academy of Science, Beijing, 100012, P.R.China", "Center for Cosmology and AstroParticle Physics, The Ohio State University, Columbus, OH 43210, USA; Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK; Lawrence Berkeley National Lab, 1 Cyclotron Rd, Berkeley, CA 94720, USA", "Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany; Universitäts-Sternwarte München, Ludwig-Maximilians-Universität München, Scheinerstraße 1, D-81679 München, Germany", "Instituto de Física Teórica, (UAM/CSIC), Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Leibniz-Institut für Astrophysik Potsdam (AIP), An der Sternwarte 16, D-14482 Potsdam, Germany", "Leibniz-Institut für Astrophysik Potsdam (AIP), An der Sternwarte 16, D-14482 Potsdam, Germany", "Instituto de Física Teórica, (UAM/CSIC), Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Campus of International Excellence UAM+CSIC, Cantoblanco, E-28049 Madrid, Spain; Departamento de Física Teórica, Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain", "Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Department of Physics and Astronomy, University of Utah, 115 S 1400 E, Salt Lake City, UT 84112, USA", "Institut de Ciències del Cosmos (ICCUB), Universitat de Barcelona (IEEC- UB), Martí i Franquès 1, E-08028 Barcelona, Spain", "Harvard-Smithsonian Center for Astrophysics, 60 Garden St, Cambridge, MA 02138, USA", "Institut Lagrange de Paris (ILP), Sorbonne Universités, 98 bis Boulevard Arago, F-75014 Paris, France; Laboratoire de Physique Nucléaire et de Hautes Energies, Université Pierre et Marie Curie, 4 Place Jussieu, F-75005 Paris, France; Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Laboratoire d'Astrophysique, Ecole Polytechnique Fédérale de Lausanne (EPFL), Observatoire de Sauverny, CH-1290 Versoix, Switzerland", "Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Department of Chemistry and Physics, King's College, 133 North River St, Wilkes Barre, PA 18711, USA", "Instituto de Física Teórica, (UAM/CSIC), Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain", "Department of Astronomy and Space Science, Sejong University, Seoul 143-747, Korea", "Universitäts-Sternwarte München, Ludwig-Maximilians-Universität Munchen, Scheinerstraße 1, D-81679 München, Germany; Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany", "Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany", "Institute of Cosmology and Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Center for Cosmology and Particle Physics, Department of Physics, New York University, 4 Washington Place, New York, NY 10003, USA", "School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews KY16 9SS, UK", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA; Center for Cosmology and AstroParticle Physics, The Ohio State University, Columbus, OH 43210, USA", "Department of Physics, Yale University, New Haven, CT 06511, USA" ]
[ "2017ApJ...849...84W", "2017JCAP...09..012O", "2017JCAP...10..009H", "2017MNRAS.464.3121W", "2017MNRAS.468.2938S", "2017MNRAS.468.4116P", "2017MNRAS.469.3762W", "2017MNRAS.470.2617A", "2017MNRAS.471.1788C", "2017MNRAS.471.2370C", "2017NatAs...1..627Z", "2017RAA....17...50Z", "2018ApJ...855...89X", "2018ApJ...857....9D", "2018ApJ...864...91S", "2018ApJ...869...26W", "2018ApJ...869L...8W", "2018JCAP...02..039L", "2018JCAP...05..033H", "2018JCAP...10..015H", "2018MNRAS.473.4773A", "2018MNRAS.477.1528W", "2018MNRAS.481.3160W", "2018PhRvD..97d3502Z", "2018PhRvD..97h3508C", "2019ApJ...877...32L", "2019ApJ...878..137Z", "2019ApJ...883..203G", "2019EPJC...79..177C", "2019JCAP...03..043J", "2019JCAP...09..010C", "2019JCAP...10..006X", "2019JCAP...10..039K", "2019MNRAS.482.3497Z", "2019MNRAS.483.1655Z", "2019MNRAS.484..442Z", "2019MNRAS.485..326L", "2019PhRvD..99f3537K", "2019RAA....19..152W", "2019arXiv191204560H", "2020Ap&SS.365...44X", "2020MNRAS.498.3470W", "2020PhRvD.101d3518Z", "2021A&A...647A..38B", "2021EPJST.230.2055M", "2021MNRAS.501.2862S", "2021MNRAS.504...33Z", "2021MNRAS.504.3956A", "2021MNRAS.505.2039P", "2021PDU....3200812R", "2021PhRvD.103b3526C", "2022MNRAS.516.5454R", "2022ScPA....2....1H", "2023ApJ...948....6T", "2023JCAP...02..061R", "2023RPPh...86b6901B", "2024arXiv240412140Y" ]
[ "astronomy" ]
6
[ "galaxies: distances and redshifts", "cosmological parameters", "cosmology: observations", "dark energy", "distance scale", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1970Ap&SS...7....3S", "1970ApJ...162..815P", "1979Natur.281..358A", "1987MNRAS.227....1K", "1988ApJ...325L..17P", "1988PhRvD..37.3406R", "1994ApJ...426...23F", "1996AJ....111.1748F", "1996MNRAS.282..877B", "1998AJ....116.1009R", "1998ApJ...496..605E", "1999ApJ...517..565P", "2000ApJ...538..473L", "2000MNRAS.312..257H", "2000PhRvL..85.4438A", "2001IJMPD..10..213C", "2002PhLB..545...23C", "2002PhRvD..66j3511L", "2003PhRvL..90i1301L", "2005ApJ...620..559J", "2005ApJ...633..560E", "2005MNRAS.362..505C", "2005PhLB..607...35F", "2005PhRvD..72l3515Z", "2006AJ....131.2332G", "2006IJMPD..15.2105S", "2006PASJ...58...93Y", "2007A&A...464..399H", "2007ApJ...664..660E", "2007ApJ...664..675E", "2007ApJ...665...14S", "2008ApJ...686...13S", "2008PhRvD..78h7303F", "2009PhRvD..80f3508P", "2010MNRAS.401.2148P", "2011AJ....142...72E", "2011MNRAS.416.3017B", "2011MNRAS.417.1350R", "2011MNRAS.418.1707B", "2012AJ....144..144B", "2012ApJ...756..127G", "2012MNRAS.427.1891A", "2012MNRAS.427.3435A", "2012PhR...513....1C", "2012PhRvD..86j3518P", "2012PhRvL.109q1301Z", "2013AJ....145...10D", "2013AJ....146...32S", "2013MNRAS.432.2433H", "2013PhR...530...87W", "2014A&A...568A..22B", "2014MNRAS.439.2531P", "2014MNRAS.441...24A", "2014PASJ...66R...1T", "2015ApJS..219...12A", "2015MNRAS.451..236Z", "2015MNRAS.453L..11B", "2015PhRvD..92h3532S", "2016A&A...594A...1P", "2016AJ....151...44D", "2016MNRAS.455.1553R", "2016MNRAS.456.4156K", "2016MNRAS.457.1770C", "2016MNRAS.457.4021L", "2016MNRAS.460.1173R", "2016MNRAS.460.1457S", "2016MNRAS.460.2453S", "2016MNRAS.460.3624S", "2016MNRAS.460.4210G", "2016MNRAS.461.2867Z", "2016PhRvL.116q1301K", "2017MNRAS.464.1168R", "2017MNRAS.464.1493S", "2017MNRAS.464.1640S", "2017MNRAS.464.2698R", "2017MNRAS.464.3121W", "2017MNRAS.464.3409B", "2017MNRAS.466.2242B", "2017MNRAS.467.2085G", "2017MNRAS.468.2938S", "2017MNRAS.469.1738S", "2017MNRAS.469.3762W", "2017MNRAS.470.2617A", "2017MNRAS.471.2370C", "2018LRR....21....2A" ]
[ "10.1093/mnras/stw3199", "10.48550/arXiv.1607.03153" ]
1607
1607.03153_arXiv.txt
\label{sec:intro} One of key science drivers of large spectroscopic galaxy surveys is to unveil the nature of dark energy (DE), the unknown energy component with a negative pressure to drive the accelerating expansion of the Universe \citep{Riess,Perlmutter}. The equation-of-state (EoS) function $w(z)$, which is the ratio of pressure over energy density of DE and is a function of redshift $z$ in general, is a proxy linking the nature of DE and its phenomenological features which can be probed by observations. For instance, a observational confirmation of $w=-1$ may suggest that DE is essentially vacuum energy, while a time-evolving $w$ can be a sign of new physics, \eg, dynamical dark energy scenarios \citep{quintessence1,quintessence2,phantom,quintom,kessence}, or a breakdown of general relativity on cosmological scales (see \citealt{MGreview} for a recent review of modified gravity theories). Therefore reconstructing $w(z)$ directly from data is an efficient way for DE studies \citep{DEreview,Zhao12,DErecon06}. The function $w(z)$ of the DE equation of state leaves imprints on the cosmic background expansion history, which can be probed by the effect of baryon acoustic oscillations (BAO) measured from galaxy surveys \citep{Eisenstein05,Cole2005}, besides other probes including supernovae Type Ia (SN Ia) \citep{Riess,Perlmutter}, cosmic microwave background (CMB) \citep{planck15} and so forth. BAO is a characteristic three-dimensional clustering pattern of galaxies at about 150 Mpc on the comoving scale, due to sound waves generated by the photon-baryon coupling in the early universe \citep{PY70,SZ70,EH98}. The BAO distance is traditionally measured using two-point correlation functions or power spectrum of galaxies. Recent studies find that higher-order statistics of galaxies \citep{BAO3pt}, or two-point clustering of voids can also be used for BAO measurements \citep{voidBAO}\footnote{In this work, we focus on galaxies as cosmic tracers thus will only refer to galaxies when discussing BAO measurements.}. Since the BAO scale is sensitive to cosmic geometry and it is largely immune to systematics \citep{imsys}, BAO is widely used as the `standard ruler' to calibrate the expansion rate of the Universe. Under assumptions that the BAO scale is the same in all directions with respect to the line-of-sight (l.o.s.) of the observer, one can probe the isotropic, one-dimensional (1D) BAO scale $D_V(z)\equiv \left[ cz (1+z)^2 D_A(z)^2 H^{-1}(z) \right]^{1/3}$, where $D_A(z)$ and $H(z)$ are the angular diameter distance and Hubble parameter at an effective redshift $z$ of the galaxy sample, using the monopole of the correlation function, or power spectrum of galaxies in redshift space. In fact, $D_A(z)$ and $H(z)$ can be separately measured when higher-order multipoles, \eg, the quadrupole and hexadecapole, are included in the analysis. This is due to the Alcock-Paczynski (AP) effect \citep{AP}: if one uses a wrong cosmology to convert redshifts into distances for the clustering analysis, the scales along and cross the l.o.s. will be dilated differently, which produces a measurable effect to break the degeneracy between $D_A$ and $H$ in the anisotropic, two-dimensional (2D) BAO analysis. The 2D BAO distances are more challenging to measure, but it is much more informative for DE studies because $w(z)$ is closely related to the first derivative of $H(z)$. The 1D and 2D BAO signals have been detected by a number of large galaxy surveys including the Sloan Digital Sky Survey (SDSS) \citep{Eisenstein05,Percival10, DR9,BAODR11,Gil-Mar,Cuesta,Acacia,Beutler16a,Beutler16b,Ross16}, the 2-degree Field Galaxy Redshift Survey (2dFGRS) \citep{Cole2005}, WiggleZ \citep{wigglez,2012arXiv1210.2130P}, the 6-degree Field Galaxy Survey (6dFGS) \citep{6dF} and so on. The Baryon Oscillation Spectroscopic Survey (BOSS) \citep{Dawson12}, part of SDSS III project \citep{SDSS3}, has reached percent level BAO measurements at $z_{\rm eff}=0.32$ and $z_{\rm eff}=0.57$ \citep{BAODR11,Gil-Mar,Cuesta,Beutler16a,Ross16} using the `low-redshift' (LOWZ; $0.15<z<0.43$) and `constant stellar mass' samples (CMASS; $0.43<z<0.7$) of Data Release (DR) 12 \citep{Alam} \footnote{The DR12 dataset is publicly available at \url{http://www.sdss.org/dr12/}}. It is true that using galaxies across wide redshift ranges can yield a precise BAO measurement at a single effective redshift, but this does not capture the tomographic information in redshift, which is required for the study of $w(z)$. Subdividing the galaxy sample into a small number of independent redshift slices and perform the BAO analysis in each slice can in principle recover the temporal information to some extent (see \citealt{Acacia} and \citealt{CF_4bins} for a three-bin and four-bin BAO analysis of the BOSS DR12 sample respectively). However, as the slice number increases, galaxies in each slice decrease, and we are at a risk of ending up with a seriously biased measurement due to large systematic uncertainties. One possible solution is to perform the BAO analysis in overlapping redshift slices. This on one hand guarantees the sufficiency of galaxy numbers in each subsample, on the other hand, it allows for a higher temporal resolution. In this work, we perform such a tomographic BAO analysis in Fourier space using the DR12 galaxy sample, and quantify the gain in dark energy studies. This paper is structured as follows. In Section 2, we describe the BOSS DR12 galaxy catalogues used for our analysis, and in Section 3, we perform a Fisher matrix forecast on this sample to determine the redshift binning, and present the power spectrum measurements. We perform the BAO analysis in Section 4, and apply our measurement to dark energy studies in Section 5, before we conclude in Section 6.
\label{sec:conclusion} The physics of baryonic acoustic oscillations has been well established to be a robust tool for cosmological studies. Specifically, the BAO measurements make it possible to reconstruct the history of the cosmic expansion, which is key to revealing the physics of the accelerating expansion of the Universe, and the nature of dark energy. Obtaining BAO measurements at as many redshifts as possible is ideal for tracing the cosmic expansion history. However, extracting the time evolution of the BAO signal is technically challenging. Na\"{\i}vely subdividing the galaxies into multiple independent redshift slices and performing BAO measurements in each slice is a straightforward solution, but the number of slices has to be limited to a small number, otherwise each individual slice would contain too few galaxies to enable a robust BAO measurement due to the low signal-to-noise ratio and issues of systematics. In this work, we solve this problem using multiple {\it overlapping} redshift slices, which allows for extracting the redshift information of the BAO signal in a large number of redshift slices. We exploit the completed DR12 combined galaxy sample of the BOSS survey, and obtain tomographic BAO measurements in nine overlapping redshift slices using the pre-reconstructed galaxy power spectrum multipoles up to the hexadecapole, after validating our data analysis pipeline using the MD-Patchy mock galaxy catalogues. Our measurement and likelihood routines compatible with CosmoMC are publicly available. We compare our measurement to that in a companion paper \citep{tomoBAO-xi}, which performs similar analysis using galaxy correlation functions derived from the same data sample, and find consistent results. For a further comparison, we derive a three-bin BAO measurement by coherently combining our tomographic measurements, and then compare to the BAO measurement presented in another companion paper \citep{Acacia}, and find an agreement \footnote{Note that, besides the different redshift binning scheme from that used in \citet{Acacia}, this work differs in two aspects: we use the fourth-order B-spline to obtain the overdensity field on the grid, which largely removes the aliasing effect; and include the hexadecapole in the BAO analysis, which we find indeed helps with the BAO constraint.}. The BAO measurements including the full covariance matrices presented in this work and a {\tt CosmoMC} patch is available at \url{https://sdss3.org//science/boss_publications.php}. We use our BAO measurements to constrain dark energy equation-of-state parameters, and find that for the CPL parametrisation, the $\Lambda$CDM model is favoured by a joint dataset of CMB, supernovae, BAO and weak lensing measurement. A more generic approach for dark energy studies using our measurement will be explored in a separate publication \citep{Zhao16}. For the BOSS DR12 sensitivity, we have seen that the dark energy FoM can differ by as much as 29\% between cases using tomographic, and non-tomographic BAO measurements. The ongoing and upcoming galaxy redshift surveys, including the eBOSS \footnote{\url{http://www.sdss.org/surveys/eboss/}} \citep{Overview}, DESI \footnote{\url{http://desi.lbl.gov/}}, Euclid \footnote{\url{http://www.euclid-ec.org/}} \citep{Euclid16}, PFS \footnote{\url{http://sumire.ipmu.jp/pfs/}} \citep{PFS}, and so on, cover a larger and larger cosmic volume, thus there is rich tomographic information in redshifts to be exploited. Besides the method developed in this work, alternatives such as the optimal redshift weighting method \citep{zweight1,zweight2,zweight3}, are being developed and applied to galaxy surveys.
16
7
1607.03153
We perform a tomographic baryon acoustic oscillations (BAO) analysis using the monopole, quadrupole and hexadecapole of the redshift-space galaxy power spectrum measured from the pre-reconstructed combined galaxy sample of the completed Sloan Digital Sky Survey Baryon Oscillation Spectroscopic Survey (BOSS) Data Release12 covering the redshift range of 0.20 &lt; z &lt; 0.75. By allowing for overlap between neighbouring redshift slices, we successfully obtained the isotropic and anisotropic BAO distance measurements within nine redshift slices to a precision of 1.5-3.4 per cent for D<SUB>V</SUB>/r<SUB>d</SUB>, 1.8-4.2 per cent for D<SUB>A</SUB>/r<SUB>d</SUB> and 3.7-7.5 per cent for H r<SUB>d</SUB>, depending on effective redshifts. We provide our BAO measurement of D<SUB>A</SUB>/r<SUB>d</SUB> and H r<SUB>d</SUB> with the full covariance matrix, which can be used for cosmological implications. Our measurements are consistent with those presented in Alam et al., in which the BAO distances are measured at three effective redshifts. We constrain dark energy parameters using our measurements and find an improvement of the Figure-of-Merit of dark energy in general due to the temporal BAO information resolved. This paper is a part of a set that analyses the final galaxy clustering data set from BOSS.
false
[ "lt", "z", "effective redshifts", "neighbouring redshift slices", "z &lt", "BAO distance measurements", "SUB", "Data Release12", "the completed Sloan Digital Sky Survey Baryon Oscillation Spectroscopic Survey", "BOSS", "d</SUB", "BAO", "the pre-reconstructed combined galaxy sample", "cosmological implications", "dark energy parameters", "the redshift range", "nine redshift slices", "three effective redshifts", "dark energy", "the redshift-space galaxy power spectrum" ]
12.176062
2.486879
161
2555647
[ "Heald, G.", "de Blok, W. J. G.", "Lucero, D.", "Carignan, C.", "Jarrett, T.", "Elson, E.", "Oozeer, N.", "Randriamampandry, T. H.", "van Zee, L." ]
2016MNRAS.462.1238H
[ "Neutral hydrogen and magnetic fields in M83 observed with the SKA Pathfinder KAT-7" ]
31
[ "CSIRO Astronomy and Space Science, 26 Dick Perry Avenue, Kensington, WA 6151, Australia; ASTRON, the Netherlands Institute for Radio Astronomy, Postbus 2, NL-7990 AA Dwingeloo, the Netherlands; Kapteyn Astronomical Institute, University of Groningen, PO Box 800, NL-9700 AV Groningen, the Netherlands", "ASTRON, the Netherlands Institute for Radio Astronomy, Postbus 2, NL-7990 AA Dwingeloo, the Netherlands; Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch 7701, South Africa; Kapteyn Astronomical Institute, University of Groningen, PO Box 800, NL-9700 AV Groningen, the Netherlands", "Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch 7701, South Africa", "Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch 7701, South Africa; Observatoire d'Astrophysique de l'Université de Ouagadougou, BP 7021, Ouagadougou 03, Burkina Faso", "Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch 7701, South Africa", "Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch 7701, South Africa", "SKA South Africa, The Park, Park Road, Pinelands, Cape Town 7405, South Africa; African Institute for Mathematical Sciences, 6-8 Melrose Road, Muizenberg 7945, South Africa; Centre for Space Research, North-West University, Potchefstroom 2520, South Africa", "Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch 7701, South Africa", "Astronomy Department, Indiana University, IN 47405, USA" ]
[ "2013pss5.book..641B", "2016mks..confE...7D", "2017A&A...603A.121C", "2017ASSL..434..209B", "2017ApJ...836..182J", "2017ApJ...839..118V", "2017ApJ...845..155B", "2017ApJ...849...51B", "2017arXiv170309345V", "2018AJ....155..233I", "2018MNRAS.477.3539N", "2018MNRAS.478..487N", "2018PASJ...70...73H", "2019ApJS..245...25J", "2019Galax...8....4B", "2019MNRAS.487.2797E", "2020A&A...640A.109W", "2020AJ....160..264B", "2020MNRAS.491.2366B", "2020MNRAS.499..379V", "2020PASA...37...17W", "2021A&C....3700502J", "2021ApJ...910...69S", "2022A&A...660A..77D", "2022A&A...665A..64W", "2023A&A...673A.146S", "2023A&A...675A..37E", "2023ApJ...947...21G", "2024A&A...681A..76R", "2024MNRAS.528.1630S", "2024arXiv240401774D" ]
[ "astronomy" ]
6
[ "galaxies: evolution", "galaxies: individual: M83", "galaxies: ISM", "galaxies: kinematics and dynamics", "galaxies: magnetic fields", "Astrophysics - Astrophysics of Galaxies" ]
[ "1981A&A...100...72H", "1984A&A...137..159S", "1988ngc..book.....T", "1989Natur.340..537S", "1991ApJ...368...60P", "1993A&A...274..687N", "1993ApJ...414...41M", "1995ASPC...77..433S", "1995PhDT.......238B", "1996MNRAS.281...27P", "1997PASA...14...52M", "1998A&AS..127..409K", "1998AJ....115.1693C", "1998MNRAS.299..189S", "2001ASPC..230..109P", "2003PASJ...55..351T", "2004AJ....128...16K", "2005A&A...441.1217B", "2005ApJ...619L..79T", "2006PASP..118..517B", "2007A&A...470..539B", "2007AJ....133..504K", "2007ASPC..376..127M", "2007ApJS..173..538T", "2008A&ARv..15..189S", "2008AIPC.1035..265V", "2008AJ....136.2563W", "2008AJ....136.2648D", "2008ExA....22..151J", "2009A&A...503..409H", "2009ARA&A..47..159B", "2009ApJ...693.1392S", "2009IEEEP..97.1522J", "2010A&A...514A..42B", "2010ApJ...720L..31B", "2011A&A...526A.118H", "2011A&A...534A.109P", "2011BASI...39..289G", "2011IJMPD..20..989N", "2011MNRAS.410.2057B", "2011MNRAS.410.2217W", "2012AJ....144...68J", "2012ARA&A..50..491P", "2012AfrSk..16..101B", "2013A&A...553A.116V", "2013AJ....145....6J", "2013AJ....146...48C", "2013ApJ...772L..28B", "2013ApJS..206...16P", "2014A&A...568A.104E", "2014A&A...570A..13M", "2014ApJ...782...90C", "2014ApJ...789..126B", "2015A&A...575A.118O", "2015A&ARv..24....4B", "2015AJ....149...54T", "2015AJ....150...47I", "2015MNRAS.450.3935L", "2015MNRAS.452.1617H", "2016A&A...585A..21F", "2016A&A...587L...3C", "2016ASPC..502...55C", "2016MNRAS.458.4210H", "2016MNRAS.460.1664F" ]
[ "10.1093/mnras/stw1698", "10.48550/arXiv.1607.03365" ]
1607
1607.03365_arXiv.txt
The evolution of galaxies is strongly influenced by their connection to the surrounding environment. Whether the dominant effect is in the form of gas accretion \citep{sancisi_etal_2008}, dynamical interactions \citep[e.g.,][]{boselli_gavazzi_2006}, or mergers \citep[e.g.,][]{blanton_moustakas_2009}, it is in the outskirts of galaxies where these evolutionary effects are clearly visible and traceable. Investigation of these effects is key to understanding the mechanisms by which galaxies continue to grow, or cease growing \citep[see for example][]{putman_etal_2012}. In the outer parts of galaxies, neutral hydrogen (\HI) has long been recognized as an excellent tracer of the effects listed above. As a necessary ingredient for future star formation, such material also represents an indication for the potential for continued evolution. By contrast, an ingredient of galactic outskirts whose role is not currently understood is the magnetic field content. Magnetic fields associated with the star-forming disks of galaxies are fairly well understood \citep[e.g.,][]{beck_2016}. However, the properties of magnetic fields in the outer parts of galaxies can so far only be inferred from large-scale synchrotron polarization patterns \citep{braun_etal_2010}, or from polarization properties of distant background radio galaxies \citep{bernet_etal_2013}. Whether the magnetic content of galaxy outskirts is important dynamically \citep[][and references therein]{benjamin_2000,elstner_etal_2014,henriksen_irwin_2016} or energetically \citep[cf.][]{beck_2007} remains unclear. Probing the outer parts of galaxies can be a difficult observational endeavour because the mass distribution is often extended over a large area. Careful observations with very high sensitivity to low surface brightness features are typically required. While great strides in this direction have been made during recent years, future radio telescopes will provide new capabilities to facilitate this line of research, both for \HI\ and magnetism. Amongst the telescopes that are anticipated in the next few years, two that will provide a large instantaneous field of view (FoV) are the Australian Square Kilometre Array Pathfinder \citep[ASKAP;][]{johnston_etal_2008} and the Aperture Tile in Focus \citep[APERTIF;][]{verheijen_etal_2008} upgrade to the Westerbork Synthesis Radio Telescope (WSRT). A different approach to bolster the speed needed to complete a large survey has been adopted by the Karoo Array Telescope \citep[MeerKAT;][]{jonas_2009}. This latter telescope, despite its comparatively small instantaneous FoV, will nevertheless be particularly well suited for high sensitivity (including low column density) observations over large areas on the sky. For this reason, several MeerKAT surveys are planned that typically aim to probe samples of very faint or distant galaxies \citep[see][]{booth_jonas_2012}. The prospect of combining MeerKAT with the Five hundred meter Aperture Spherical radio Telescope \citep[FAST;][]{nan_etal_2011} holds the promise of extremely sensitive observations that may probe the galaxy-IGM interface \citep{carignan_2016}. One of the several forthcoming MeerKAT surveys that are currently planned is named ``MeerKAT \HI\ Observations of Nearby Galactic Objects: Observing Southern Emitters'' (MHONGOOSE)\footnote{PI W.~J.~G.~de~Blok; see \url{http://mhongoose.astron.nl/}}. The motivation and design of MHONGOOSE is similar in many respects to the recent Westerbork Synthesis Radio Telescope (WSRT) Hydrogen Accretion in LOcal GAlaxieS \citep[HALOGAS;][]{heald_etal_2011} survey. Specifically, MHONGOOSE will perform deep \HI\ observations of thirty (30) nearby galaxies, achieving column density sensitivities of several $\times10^{18}\,\mathrm{atoms\,cm^{-2}}$ at kpc-scale physical resolution. MHONGHOOSE will build on the capability of most existing \HI\ surveys by also providing excellent polarization data along with the spectral line cubes, thus enabling investigation of the magnetic fields along with the gas morphology and kinematics. As the community prepares for data and science results from MHONGOOSE, initial exploratory observations are being performed with the KAT-7 radio telescope \citep[see, e.g.,][]{lucero_etal_2015,hess_etal_2015,carignan_etal_2016}, which is the precursor radio telescope located on the MeerKAT site in South Africa \citep{carignan_etal_2013}. M~83 (NGC~5236) is a grand-design spiral galaxy at the center of a loose group \citep{karachentsev_etal_2007}. M~83 is a prominent example of an interesting class of galaxies showing extended UV (XUV) emission \citep{thilker_etal_2005,thilker_etal_2007}, indicative of active star formation taking place out to several optical radii (nearly $4\times\,r_{25}$). What fuels this outer disk star formation? Recent \HI\ observations of M83 have been obtained with the Australia Telescope Compact Array \citep[ATCA;][]{park_etal_2001,jarrett_etal_2013} and with the Very Large Array \citep[VLA;][]{walter_etal_2008,deblok_etal_2008,barnes_etal_2014}. These reveal a very extended \HI\ distribution that also reaches far beyond the main optical disk, and with a very close morphological correlation between sites of recent star formation and high-column density regions in the \HI\ reservoir \citep{bigiel_etal_2010}. With the available data \citet{bigiel_etal_2010} were able to constrain the outer-disk gas depletion time from {\it in situ} star formation to of order 100~Gyr. In this paper we present new \HI\ observations that reveal even more neutral gas mass in the outer disk than has been seen previously, strengthening the conclusion that a vast, nearly untapped gas mass resides outside of the main disk in this galaxy, and may provide a source of fuel to maintain star formation occurring within the central star forming disk. We summarize some basic properties of M~83 in Table~\ref{table:m83properties}. This paper is organized as follows. We describe the KAT-7 data collection and reduction procedures in \S\,\ref{section:data}. The data are described in \S\,\ref{section:analysis}, where we present new details about the \HI\ distribution and kinematics, along with their connection to stellar features in the outer disk and to the environment more generally. We also present new conclusions regarding the large-scale magnetic field structure in M~83 (\S\,\ref{section:polarization}). We conclude the paper in \S\,\ref{section:conclusions}. \begin{table*} \centering \begin{minipage}{140mm} \caption{Properties of M83.} \begin{tabular}{@{}lll@{}} \hline Property & Value & Reference \\ \hline Hubble type & SAB(s)c & \citet{devaucouleurs_etal_1991}\\ Distance & 4.79~Mpc ($43.1\arcsec=1\,\mathrm{kpc}$) & \citet{karachentsev_etal_2007}\\ $D_{25}$ & $11\farcm7$ & \citet{tully_1988}\\ $M_B$ & -20.94 & \citet{makarov_etal_2014}\\ SFR (12$\mu$m) & $4.95\pm0.09\,M_\odot\,\mathrm{yr}^{-1}$ & \citet{jarrett_etal_2013,cluver_etal_2014}\\ SFR (22$\mu$m) & $3.86\pm0.07\,M_\odot\,\mathrm{yr}^{-1}$ & \citet{jarrett_etal_2013,cluver_etal_2014}\\ \hline Total \HI\ mass & $9.0\times10^9\,M_\odot$ & This work\\ Stellar mass & $2.88\times10^{10}\,M_\odot$ & \citet{jarrett_etal_2013,cluver_etal_2014}\\ Virial mass & $2\times10^{12}\,M_\odot$ & \citet{tully_2015} \\ Systemic velocity & $510\,\mathrm{km\,s}^{-1}$ & This work\\ Maximum rotational velocity & $170\,\mathrm{km\,s}^{-1}$ & This work\\ \hline \end{tabular}\label{table:m83properties} \end{minipage} \end{table*}
16
7
1607.03365
We present new KAT-7 observations of the neutral hydrogen (H I) spectral line, and polarized radio continuum emission, in the grand-design spiral M83. These observations provide a sensitive probe of the outer-disc structure and kinematics, revealing a vast and massive neutral gas distribution that appears to be tightly coupled to the interaction of the galaxy with the environment. We present a new rotation curve extending out to a radius of 50 kpc. Based on our new H I data set and comparison with multiwavelength data from the literature, we consider the impact of mergers on the outer disc and discuss the evolution of M83. We also study the periphery of the H I distribution and reveal a sharp edge to the gaseous disc that is consistent with photoionization or ram pressure from the intergalactic medium. The radio continuum emission is not nearly as extended as the H I and is restricted to the main optical disc. Despite the relatively low angular resolution, we are able to draw broad conclusions about the large-scale magnetic field topology. We show that the magnetic field of M83 is similar in form to other nearby star-forming galaxies, and suggest that the disc-halo interface may host a large-scale regular magnetic field.
false
[ "M83", "other nearby star-forming galaxies", "form", "polarized radio continuum emission", "the main optical disc", "the large-scale magnetic field topology", "a large-scale regular magnetic field", "the outer disc", "the gaseous disc", "spectral line", "the magnetic field", "the disc-halo interface", "pressure", "broad conclusions", "multiwavelength data", "photoionization", "the outer-disc structure", "kinematics", "the grand-design spiral M83", "the intergalactic medium" ]
12.595153
9.091754
-1
12623388
[ "Brandenburg, Axel", "Kahniashvili, Tina" ]
2017PhRvL.118e5102B
[ "Classes of Hydrodynamic and Magnetohydrodynamic Turbulent Decay" ]
69
[ "Laboratory for Atmospheric and Space Physics, University of Colorado, Boulder, Colorado 80303, USA; JILA and Department of Astrophysical and Planetary Sciences, University of Colorado, Boulder, Colorado 80303, USA; Nordita, KTH Royal Institute of Technology and Stockholm University, Roslagstullsbacken 23, 10691 Stockholm, Sweden; Department of Astronomy, AlbaNova University Center, Stockholm University, 10691 Stockholm, Sweden", "The McWilliams Center for Cosmology and Department of Physics, Carnegie Mellon University, 5000 Forbes Avenue, Pittsburgh, Pennsylvania 15213, USA; Department of Physics, Laurentian University, Ramsey Lake Road, Sudbury, ON P3E 2C, Canada; Abastumani Astrophysical Observatory, Ilia State University, 3-5 Cholokashvili Ave, Tbilisi, GE-0194, Georgia" ]
[ "2016PhRvD..94j3510C", "2016PhyS...91j4008K", "2017ApJ...845L..21B", "2017ApJ...850L...8S", "2017JCAP...12..002K", "2017MNRAS.472.1628P", "2017PhRvD..96h3505P", "2017PhRvD..96h3511S", "2017PhRvD..96l3528B", "2018AN....339..641B", "2018ApJ...858..124S", "2018CQGra..35l4003C", "2018CQGra..35l4004P", "2018CQGra..35n4001P", "2018JCAP...08..034B", "2018JPlPh..84d7304B", "2018PhRvD..97h3503S", "2019ApJ...870...87B", "2019JCAP...11..028P", "2019MNRAS.484..185P", "2019PhRvE..99a3101M", "2019PhRvF...4b4608B", "2020ApJ...889...55B", "2020ApJ...901...18B", "2020GApFD.114..130R", "2020IAUGA..30..295K", "2020PhFl...32i5109A", "2020PhRvD.101j3526S", "2020PhRvD.102b3536B", "2021ApJ...911..110B", "2021ApJ...920...26B", "2021ApJ...922..192B", "2021CQGra..38n5002B", "2021JCAP...04..034K", "2021JFM...916A...4Y", "2021JPlPh..87f9020Z", "2021MNRAS.500.5350V", "2021PhRvD.104d3513B", "2021PhRvX..11d1005H", "2021RPPh...84g4901V", "2022ApJ...929..127M", "2022JCAP...04..019R", "2022JCAP...09..029A", "2022JPlPh..88e1501S", "2022JPlPh..88f9002Z", "2022MNRAS.517.3916P", "2022PhRvD.105d1302S", "2022PhRvD.105l3502R", "2022PhRvD.106f3511D", "2022PhRvD.106j3536S", "2023ApJ...944....2P", "2023Atmos..14..932B", "2023Entrp..25.1270B", "2023JFM...973A..13H", "2023JHEP...01..053D", "2023JPlPh..89a1701B", "2023JPlPh..89f9006B", "2023MNRAS.518.3312B", "2023MNRAS.520.6268H", "2023PDU....4001212C", "2023PhLB..84338002U", "2023PhRvD.107j3524H", "2023PhRvD.108f3529B", "2023PhRvE.107e5206A", "2023PhRvR...5b2028B", "2023PrPNP.12904016K", "2024arXiv240108569B", "2024arXiv240316763P", "2024arXiv240611798B" ]
[ "astronomy", "physics" ]
14
[ "Physics - Fluid Dynamics", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1956RSPTA.248..369B", "1967PhFl...10.1349S", "1976JFM....77..321P", "1979JFM....93..631C", "1983JAtS...40..749L", "1992PhFlA...4.1492G", "1996PhRvD..54.1291B", "1997PhLB..398..321O", "1998PhRvL..80.2754M", "1999PhRvL..82.4831S", "1999PhRvL..83.2195B", "2001PhRvE..64e6405C", "2002CoPhC.147..471B", "2003JFM...480..129K", "2004PhRvD..70h3009C", "2004PhRvD..70l3003B", "2004PhRvE..69d6304K", "2004PhRvE..69e6303Y", "2004PhRvE..70b6405H", "2006MNRAS.366.1437S", "2010JFM...663..268D", "2012ApJ...759...54T", "2013PhRvD..87h3007K", "2014ApJ...794L..26Z", "2015JPlPh..81e3401F", "2015PhRvL.114g5001B", "2015arXiv150908962O", "2015arXiv151105007O", "2016EPJC...76..504C", "2016JCAP...01..002W" ]
[ "10.1103/PhysRevLett.118.055102", "10.48550/arXiv.1607.01360" ]
1607
1607.01360_arXiv.txt
By default, the {\sc Pencil Code} uses sixth order accurate finite difference representations for the first and second derivatives. A low spatial order of the scheme implies that at high wavenumbers the magnitude of the numerical derivative is reduced, leading to lower advection speeds of the high wavenumber Fourier components. This is generally referred to as phase error. Thus, for an advected tophat function, the high wavenumber constituents will lag behind, creating the well-known Gibbs phenomenon which needs to be controlled by a certain amount of viscosity. Higher order schemes require less viscosity to control the Gibbs phenomenon \citep{BD02}. On the other hand, any turbulence simulation requires a sufficient amount of viscosity to dissipate kinetic energy. It is therefore thought that for a sixth orders scheme the two limits on the viscosity are similar and that it is not advantageous to use higher order representations of the spatial derivatives. \begin{table}[b!]\caption{ Coefficients $c_j^{(n)}\equiv a_j^{(n)}/b^{(n)}$ }\vspace{12pt}\centerline{\begin{tabular}{rcccccccc} $N$ & $n$ & $b^{(n)}$ & $a_0^{(n)}$ & $a_1^{(n)}$ & $a_2^{(n)}$ & $a_3^{(n)}$ & $a_4^{(n)}$ & $a_5^{(n)}$ \\ \hline 10 & 1 & 2520 & 0 & 2100 & $-600$ & 150 & $-25$ & 2 \\ 8 & 1 & 840 & 0 & 672 & $-168$ & 32 & $-3$ & \\ 6 & 1 & 60 & 0 & 45 & $-9$ & 1 & & \\ 4 & 1 & 12 & 0 & 8 & $-1$ & & & \\ 2 & 1 & 2 & 0 & 1 & & & & \\ 10 & 2 & 25200 & $-73766$ & 42000 & $-6000$ & $1000$ & $-125$ & $8$ \\ 8 & 2 & 5040 & $-14350$ & 8064 & $-1008$ & $128$ & $-9$ & \\ 6 & 2 & 180 & $-490$ & 270 & $-27$ & $2$ & & \\ 4 & 2 & 12 & $-30$ & 16 & $-1$ & & & \\ 2 & 2 & 1 & $-2$ & 1 & & & & \\ \label{Tcoeff}\end{tabular}}\end{table} \begin{figure}[t!]\begin{center} \includegraphics[width=.9\columnwidth]{pspec_comp_10th} \end{center}\caption[]{ Magnetic (upper curves) and kinetic (lower curves) energy spectra for at $t=110$ for the sixth order (blue, dashed) and tenth order (red, solid) finite difference schemes. }\label{pspec_comp_10th}\end{figure} To verify this in the present context, we have run a high Reynolds number case both with sixth and tenth order schemes. In the {\sc Pencil Code}, the order of the scheme can easily be changed by setting {\tt DERIV=deriv\_10th}. In that case, first and second derivatives are represented as \EQ \dd^n f_i/\dd x^n=\sum_{j=-N}^N (\sgn j)^n c_{|j|}^{(n)} f_{i+j}/\delta x^n, \EN with coefficient $c_j^{(n)}$ given in \Tab{Tcoeff} for schemes of order $N$. The result of the comparison is shown in \Fig{pspec_comp_10th}. The differences between the two cases are negligible, except that with the more accurate tenth order scheme the inverse transfer of kinetic energy to larger scales is now slightly stronger. This is consistent with our earlier findings that the inverse transfer in nonhelical MHD becomes more pronounced at larger resolution.
16
7
1607.01360
We perform numerical simulations of decaying hydrodynamic and magnetohydrodynamic turbulence. We classify our time-dependent solutions by their evolutionary tracks in parametric plots between instantaneous scaling exponents. We find distinct classes of solutions evolving along specific trajectories toward points on a line of self-similar solutions. These trajectories are determined by the underlying physics governing individual cases, while the infrared slope of the initial conditions plays only a limited role. In the helical case, even for a scale-invariant initial spectrum (inversely proportional to wave number k ), the solution evolves along the same trajectory as for a Batchelor spectrum (proportional to k<SUP>4</SUP>).
false
[ "solutions", "instantaneous scaling exponents", "specific trajectories", "k", "parametric plots", "self-similar solutions", "individual cases", "Batchelor", "number", "points", "the same trajectory", "the solution", "a Batchelor spectrum", "hydrodynamic and magnetohydrodynamic turbulence", "our time-dependent solutions", "distinct classes", "numerical simulations", "a scale-invariant initial spectrum", "the initial conditions", "These trajectories" ]
11.976563
14.060761
2
12468275
[ "Suh, Kyung-Won" ]
2016JKAS...49..127S
[ "Optical Properties of Amorphous Alumina Dust in the Envelopes around O-Rich AGB Stars" ]
7
[ "Department of Astronomy and Space Science, Chungbuk National University, Cheongju 28644, Korea" ]
[ "2018A&A...620A..75K", "2018JKAS...51..155S", "2019IAUS..343...31K", "2019IAUS..343..159S", "2020ApJ...891...43S", "2021ApJS..256...43S", "2022JKAS...55..195S" ]
[ "astronomy" ]
3
[ "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1976ApJ...209..509J", "1983asls.book.....B", "1995A&AS..112..143H", "1995Icar..114..203K", "1997ApJ...476..199B", "1997MNRAS.287..799I", "1998ApJS..119..141S", "1999MNRAS.304..389S", "2000A&AS..146..437S", "2001ApJ...558..165E", "2003yCat.2246....0C", "2003yCat.5114....0E", "2004ApJ...615..485S", "2006A&A...450.1051J", "2007ApJ...668L.107M", "2007PASJ...59S.369M", "2011MNRAS.417.3047S", "2012JKAS...45..139K", "2013A&A...553A..81Z", "2013ApJ...762..113S", "2014MNRAS.440..631J", "2015ApJ...808..165S" ]
[ "10.5303/JKAS.2016.49.4.127", "10.48550/arXiv.1607.05363" ]
1607
1607.05363_arXiv.txt
\label{introduction} The main site of dust formation is believed to be the cool envelopes around asymptotic giant branch (AGB) stars. O-rich AGB stars (M-type Miras and OH/IR stars) typically show conspicuous 10 $\mu$m and 18 $\mu$m features in emission or absorption. They suggest the presence of amorphous silicate dust grains in the outer envelopes around them (Jones \& Merrill 1976). Low mass-loss rate O-rich AGB (LMOA) stars with thin dust envelopes show the 10 $\mu$m and 18 $\mu$m emission features of amorphous silicate. High mass-loss rate O-rich AGB (HMOA) stars with thick dust envelopes show the absorbing features at the same wavelengths (e.g., Suh 2004). Water-ice was found in some HMOA stars (Justtanont et al. 2006; Suh \& Kwon 2013) and amorphous alumina (Al$_2$O$_3$) dust grains were detected in many LMOA stars (e.g., Sloan \& Price 1998). In a number of previous works (e.g., Markwick-Kemper et al. 2007; Suh \& Kwon 2011), the term corundum was erroneously assigned to different kinds of solid alumina regardless of its crystal structure. However, only one of the crystalline Al$_2$O$_3$ polymorphs deserves the name corundum: namely $\alpha$-Al$_2$O$_3$ which has trigonal lattice symmetry (e.g., Koike et al. 1995). While the amorphous aluminum oxide material (Al$_2$O$_3$) synthesized by Begemann et al. (1997) shows a single peak at 11.8 $\mu$m, corundum shows much sharper multiple peaks around 12 $\mu$m (Koike et al. 1995; Zeidler et al. 2013). In this work, we investigate optical properties of amorphous alumina (Al$_2$O$_3$) dust in the envelopes around O-rich AGB stars. We derive the optical constants of the alumina dust in a wide wavelength range, which satisfy the Kramers-Kronig relation and reproduce the laboratory measured optical data. Using the opacity function of the amorphous alumina dust, we compare the theoretical radiative transfer model results with the observed spectral energy distributions (SEDs) and observations on various IR two-color diagrams (2CDs) for a large sample O-rich AGB stars.
\label{summary} In this work, we have investigated optical properties of the amorphous alumina (Al$_2$O$_3$) dust grains in the envelopes around O-rich AGB stars, considering the laboratory measured optical data. We have derived the the optical constants of the amorphous alumina dust in a wide wavelength range, which satisfy the Kramers-Kronig relation and reproduce the laboratory measured data. Amorphous alumina grains produce a single peak at 11.8 $\mu$m and influences the shape of the SED at around 10 $\mu$m. The shape of the 10 $\mu$m feature of O-rich AGB stars, which is mainly produced by silicate, can be modified by addition of the alumina dust. Using the opacity function of the alumina dust, we have compared the theoretical radiative transfer model results with the observed SEDs and observations on various IR 2CDs for a large sample O-rich AGB stars. Even though it is difficult to suggest the exact content of amorphous alumina for O-rich AGB stars because dust species other than alumina can also produce similar features in the wavelength range 8-15 $\mu$m, we have found a general trend for a large sample of the stars on 2CDs. Comparing the theoretical models with the observations on various IR 2CDs, we have found that the amorphous alumina dust (about 10-40 \%) mixed with amorphous silicate can reproduce much more observed points for LMOA stars, which have thin dust envelopes. Because the alumina dust is not useful for HMOA stars, we expect that the relative alumina abundance for LMOA stars is higher than the abundance for HMOA stars with thick dust envelopes. We expect that the optical constants for amorphous alumina derived in this work would be useful for further studies on dust around AGB and post-AGB stars. The optical constants for the amorphous alumina derived in this work will be accessible through the author's web site: \url{http://web.chungbuk.ac.kr/~kwsuh/d-opt.htm}.
16
7
1607.05363
We investigate optical properties of amorphous alumina (Al_2O_3) dust grains in the envelopes around O-rich asymptotic giant branch (AGB) stars considering the laboratory measured optical data. We derive the optical constants of amorphous alumina in a wide wavelength range that satisfy the Kramers-Kronig relation and reproduce the laboratory measured data. Using the amorphous alumina and silicate dust, we compare the radiative transfer model results with the observed spectral energy distributions. Comparing the theoretical models with the observations on various IR two-color diagrams for a large sample O-rich AGB stars, we find that the amorphous alumina dust (about 10-40 %) mixed with amorphous silicate can reproduce much more observed points for the O-rich AGB stars with thin dust envelopes.
false
[ "measured optical data", "thin dust envelopes", "amorphous alumina", "dust", "amorphous silicate", "optical properties", "the amorphous alumina dust", "a large sample O-rich AGB stars", "AGB", "the laboratory measured data", "the O-rich AGB stars", "the observed spectral energy distributions", "the amorphous alumina", "amorphous alumina (Al_2O_3) dust grains", "Kramers", "the radiative transfer model results", "IR", "Al_2O_3", "a wide wavelength range", "the optical constants" ]
9.732187
11.229342
123
12520735
[ "Agostini, M.", "Altenmüller, K.", "Appel, S.", "Atroshchenko, V.", "Bellini, G.", "Benziger, J.", "Bick, D.", "Bonfini, G.", "Bravo, D.", "Caccianiga, B.", "Calaprice, F.", "Caminata, A.", "Carlini, M.", "Cavalcante, P.", "Chepurnov, A.", "Choi, K.", "D'Angelo, D.", "Davini, S.", "de Kerret, H.", "Derbin, A.", "Di Noto, L.", "Drachnev, I.", "Etenko, A.", "Fomenko, K.", "Franco, D.", "Gabriele, F.", "Galbiati, C.", "Ghiano, C.", "Giammarchi, M.", "Goeger-Neff, M.", "Goretti, A.", "Gromov, M.", "Hagner, C.", "Hungerford, E.", "Ianni, Aldo", "Ianni, Andrea", "Jany, A.", "Jedrzejczak, K.", "Jeschke, D.", "Kobychev, V.", "Korablev, D.", "Korga, G.", "Kryn, D.", "Laubenstein, M.", "Lehnert, B.", "Litvinovich, E.", "Lombardi, F.", "Lombardi, P.", "Ludhova, L.", "Lukyanchenko, G.", "Machulin, I.", "Manecki, S.", "Maneschg, W.", "Manuzio, G.", "Marcocci, S.", "Meroni, E.", "Meyer, M.", "Miramonti, L.", "Misiaszek, M.", "Montuschi, M.", "Mosteiro, P.", "Muratova, V.", "Neumair, B.", "Oberauer, L.", "Obolensky, M.", "Ortica, F.", "Pallavicini, M.", "Papp, L.", "Pocar, A.", "Ranucci, G.", "Razeto, A.", "Re, A.", "Romani, A.", "Roncin, R.", "Rossi, N.", "Schönert, S.", "Semenov, D.", "Skorokhvatov, M.", "Smirnov, O.", "Sotnikov, A.", "Sukhotin, S.", "Suvorov, Y.", "Tartaglia, R.", "Testera, G.", "Thurn, J.", "Toropova, M.", "Unzhakov, E.", "Vishneva, A.", "Vogelaar, R. B.", "von Feilitzsch, F.", "Wang, H.", "Weinz, S.", "Winter, J.", "Wojcik, M.", "Wurm, M.", "Yokley, Z.", "Zaimidoroga, O.", "Zavatarelli, S.", "Zuber, K.", "Zuzel, G.", "Borexino Collaboration" ]
2017APh....86...11A
[ "Borexino's search for low-energy neutrino and antineutrino signals correlated with gamma-ray bursts" ]
13
[ "Gran Sasso Science Institute (INFN), 67100 Ł'Aquila, Italy", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "Chemical Engineering Department, Princeton University, Princeton, NJ 08544, USA", "Institut für Experimentalphysik, Universität, 22761 Hamburg, Germany", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Physics Department, Virginia Polytechnic Institute and State University, Blacksburg, VA 24061, USA", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "Physics Department, Princeton University, Princeton, NJ 08544, USA", "Dipartimento di Fisica, Università degli Studi e INFN, Genova 16146, Italy", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy; Physics Department, Virginia Polytechnic Institute and State University, Blacksburg, VA 24061, USA", "Lomonosov Moscow State University Skobeltsyn Institute of Nuclear Physics, 119234 Moscow, Russia", "Department of Physics and Astronomy, University of Hawaii, Honolulu, HI 96822, USA", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "Gran Sasso Science Institute (INFN), 67100 Ł'Aquila, Italy", "AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/IRFU, Observatoire de Paris, Sorbonne Paris Cité, 75205 Paris Cedex 13, France", "St. Petersburg Nuclear Physics Institute NRC Kurchatov Institute, 188350 Gatchina, Russia", "Dipartimento di Fisica, Università degli Studi e INFN, Genova 16146, Italy", "Gran Sasso Science Institute (INFN), 67100 Ł'Aquila, Italy", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia", "Joint Institute for Nuclear Research, 141980 Dubna, Russia", "AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/IRFU, Observatoire de Paris, Sorbonne Paris Cité, 75205 Paris Cedex 13, France", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Physics Department, Princeton University, Princeton, NJ 08544, USA", "Dipartimento di Fisica, Università degli Studi e INFN, Genova 16146, Italy", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "Physics Department, Princeton University, Princeton, NJ 08544, USA", "Lomonosov Moscow State University Skobeltsyn Institute of Nuclear Physics, 119234 Moscow, Russia", "Institut für Experimentalphysik, Universität, 22761 Hamburg, Germany", "Department of Physics, University of Houston, Houston, TX 77204, USA", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy; Laboratorio Subterráneo de Canfranc, Paseo de los Ayerbe S/N, 22880 Canfranc Estacion Huesca, Spain", "Physics Department, Princeton University, Princeton, NJ 08544, USA", "M. Smoluchowski Institute of Physics, Jagiellonian University, 30059 Krakow, Poland", "M. Smoluchowski Institute of Physics, Jagiellonian University, 30059 Krakow, Poland", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "Institute for Nuclear Research, 03680 Kiev, Ukraine", "Joint Institute for Nuclear Research, 141980 Dubna, Russia", "Department of Physics, University of Houston, Houston, TX 77204, USA", "AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/IRFU, Observatoire de Paris, Sorbonne Paris Cité, 75205 Paris Cedex 13, France", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Department of Physics, Technische Universität Dresden, 01062 Dresden, Germany", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia; National Research Nuclear University MEPhI (Moscow Engineering Physics Institute), 115409 Moscow, Russia", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "IKP-2 Forschungzentrum Jülich, 52428 Jülich, Germany; RWTH Aachen University, 52062 Aachen, Germany", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia; National Research Nuclear University MEPhI (Moscow Engineering Physics Institute), 115409 Moscow, Russia", "Physics Department, Virginia Polytechnic Institute and State University, Blacksburg, VA 24061, USA; Physics Department, Queen's University, Kingston ON K7L 3N6, Canada", "Max-Planck-Institut für Kernphysik, 69117 Heidelberg, Germany", "Dipartimento di Fisica, Università degli Studi e INFN, Genova 16146, Italy", "Gran Sasso Science Institute (INFN), 67100 Ł'Aquila, Italy", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "Institut für Experimentalphysik, Universität, 22761 Hamburg, Germany", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "M. Smoluchowski Institute of Physics, Jagiellonian University, 30059 Krakow, Poland; INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Dipartimento di Fisica e Scienze della Terra, Università degli Studi di Ferrara e INFN, Via Saragat 1-44122, Ferrara, Italy", "Physics Department, Princeton University, Princeton, NJ 08544, USA", "St. Petersburg Nuclear Physics Institute NRC Kurchatov Institute, 188350 Gatchina, Russia", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/IRFU, Observatoire de Paris, Sorbonne Paris Cité, 75205 Paris Cedex 13, France", "Dipartimento di Chimica, Biologia e Biotecnologie, Università e INFN, 06123 Perugia, Italy", "Dipartimento di Fisica, Università degli Studi e INFN, Genova 16146, Italy", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "Amherst Center for Fundamental Interactions and Physics Department, University of Massachusetts, Amherst, MA 01003, USA", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Dipartimento di Fisica, Università degli Studi e INFN, 20133 Milano, Italy", "Dipartimento di Chimica, Biologia e Biotecnologie, Università e INFN, 06123 Perugia, Italy", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy; AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/IRFU, Observatoire de Paris, Sorbonne Paris Cité, 75205 Paris Cedex 13, France", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "St. Petersburg Nuclear Physics Institute NRC Kurchatov Institute, 188350 Gatchina, Russia", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia; National Research Nuclear University MEPhI (Moscow Engineering Physics Institute), 115409 Moscow, Russia", "Joint Institute for Nuclear Research, 141980 Dubna, Russia", "Joint Institute for Nuclear Research, 141980 Dubna, Russia", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia", "Physics and Astronomy Department, University of California Los Angeles (UCLA), Los Angeles, California 90095, USA; National Research Centre Kurchatov Institute, 123182 Moscow, Russia", "INFN Laboratori Nazionali del Gran Sasso, 67010 Assergi (AQ), Italy", "Dipartimento di Fisica, Università degli Studi e INFN, Genova 16146, Italy", "Department of Physics, Technische Universität Dresden, 01062 Dresden, Germany", "National Research Centre Kurchatov Institute, 123182 Moscow, Russia", "St. Petersburg Nuclear Physics Institute NRC Kurchatov Institute, 188350 Gatchina, Russia", "Joint Institute for Nuclear Research, 141980 Dubna, Russia", "Physics Department, Virginia Polytechnic Institute and State University, Blacksburg, VA 24061, USA", "Physik-Department and Excellence Cluster Universe, Technische Universität München, 85748 Garching, Germany", "Physics and Astronomy Department, University of California Los Angeles (UCLA), Los Angeles, California 90095, USA", "Institute of Physics and Excellence Cluster PRISMA, Johannes Gutenberg-Universität Mainz, 55099 Mainz, Germany", "Institute of Physics and Excellence Cluster PRISMA, Johannes Gutenberg-Universität Mainz, 55099 Mainz, Germany", "M. Smoluchowski Institute of Physics, Jagiellonian University, 30059 Krakow, Poland", "Institute of Physics and Excellence Cluster PRISMA, Johannes Gutenberg-Universität Mainz, 55099 Mainz, Germany", "Physics Department, Virginia Polytechnic Institute and State University, Blacksburg, VA 24061, USA", "Joint Institute for Nuclear Research, 141980 Dubna, Russia", "Dipartimento di Fisica, Università degli Studi e INFN, Genova 16146, Italy", "Department of Physics, Technische Universität Dresden, 01062 Dresden, Germany", "M. Smoluchowski Institute of Physics, Jagiellonian University, 30059 Krakow, Poland", "-" ]
[ "2017ApJ...850...21A", "2017arXiv170900756T", "2018APh....97..136A", "2018ApJ...852..120B", "2018IJMPA..3343009C", "2018JMPh....9..573T", "2018JMPh....9.1827T", "2020JPhCS1342a2035M", "2020PhRvD.101f2001A", "2021APh...12502509A", "2021PTEP.2021j3F01O", "2021Univ....7..231K", "2022ApJ...927...69A" ]
[ "astronomy" ]
4
[ "Neutrinos", "Antineutrinos", "Gamma-ray bursts", "Low energy/MeV neutrinos", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1996PhRvD..54.2779H", "1998PhRvD..57.3873F", "2001PhRvD..64d3004B", "2001PhRvL..87q1102M", "2002ApJ...578..317F", "2003PhLB..564...42S", "2005PhRvD..71d7303S", "2006ApJ...641..961L", "2006RPPh...69.2259M", "2007ApJ...664..397A", "2007RSPTA.365.1323W", "2009NIMPA.600..568A", "2009NIMPA.609...58A", "2010ApJ...710..346A", "2010PhRvL.104x1102C", "2011ApJ...736...50V", "2011PhLB..696..191B", "2013A&A...559A...9A", "2013JCAP...03..006A", "2014APh....55....1A", "2014APh....57...16A", "2014Natur.512..383B", "2014PhRvD..89k2007B", "2015ApJ...805L...5A", "2015ApJ...806...87A", "2015PhRvD..92c1101A" ]
[ "10.1016/j.astropartphys.2016.10.004", "10.48550/arXiv.1607.05649" ]
1607
1607.05649_arXiv.txt
\label{sec:intro} Gamma ray bursts (GRBs) are among the most energetic events known in the Universe, with a typical apparent energy release of $~10^{54}$ erg (or $\sim$1 solar mass), assuming isotropic emission of energy. The average rate of observed GRBs is about 1 event per day from the entire sky. The observer-frame duration of gamma ray emission in the MeV range can be less than 2\,s (for a smaller sub-class of so called short GRBs) or, more typically, is of the order of 10 to 100 seconds. Longer afterglows in X-rays, optical, and radio wavelengths are also observed. The measured redshifts of optical afterglows ($z= 0.01..8.2$) allow to attribute GRBs as extra-galactic events, whose progenitors lie at cosmological distances. Currently, there is no universally accepted model of GRBs. However, the long GRBs are usually linked to the rotating cores of very massive stars collapsed into neutron stars (NS) or black holes (BH) \cite{Lee 2006, Meszaros 2006}. Short GRBs can result from binary mergers of NS + NS or NS + BH \cite{Meszaros 2006}. These models usually assume that neutrino cooling dominates over the electromagnetic one with neutrino energies in the MeV range \cite{Sahu 2005}. In these models, the energy emitted in the form of MeV thermal neutrinos is of the order of one Solar mass ($M_\odot c^2\approx 2\cdot10^{54}$~ergs) while the energy released in gamma quanta is up to $\sim$100 times less~\cite{Halzen 1996, Meszaros 2015}. Even for the nearest GRBs with $z \sim 0.01$ (for the standard $\Lambda$CDM cosmological model~\cite{Planck}, this comoving distance is $\sim 10^{26}\, {\rm cm}$ or 30 Mpc), the expected fluence of low-energy neutrinos at Earth is equal to $10^5-10^6\, {\rm cm}^{-2}$, which is too small to be observed by current detectors. However, there exist other models for the origin of GRBs, such as cusps of superconducting cosmic strings~\cite{Berezinsky 2001} which can better explain some peculiarities of GRBs~\cite{Cheng 2010}. Some of these models are predicting the fluence of MeV-range neutrinos to be up to $10^{10}$ times larger than that of gamma quanta. The neutrino fluence in these models can be estimated as~\cite{Halzen 1996}: \begin{equation}\label{eq:eq0} \begin{split} \Phi_\nu &= 10^8\,{\rm cm}^{-2}\, \left(\frac{\eta_\gamma}{10^{-10}}\right)^{-1} \left(\frac{E_\nu}{100\,{\rm MeV}}\right)^{-1} \\ & \times\left(\frac{F_\gamma}{10^{-6}\,{\rm erg\cdot cm}^{-2}}\right), \end{split} \end{equation} where $\eta_\gamma$ is the ratio of photon and neutrino fluences, $F_\gamma$ is the observed gamma ray energy fluence of the GRB. Thus, for a typical GRB with gamma energy release of $3\times 10^{51}\,{\rm erg}$ at redshift $z = 2$ ($F_\gamma=10^{-6}\,{\rm erg\cdot cm}^{-2}$), the predicted fluence of neutrinos with energy of $\sim$10~MeV is $\sim$$10^8 -10^9\, {\rm cm}^{-2}$. For comparison, the {\it observed} flux of solar neutrinos is about $6\cdot 10^{10}\, {\rm cm}^{-2}\cdot{\rm s}^{-1}$, of geo-antineutrinos is about $5\cdot 10^6\, {\rm cm}^{-2}\cdot{\rm s}^{-1}$~\cite{BX2015geonu}. It demonstrates that the sensitivity level of existing neutrino detectors in the MeV range is close to the fluxes expected in several GRB models, if one uses a big set of GRBs. Hereinafter, the energy of neutrino refers to the observed energy; the emitted energy has to be multiplied by factor $1+z$ which is not important due to large uncertainties in models predictions. Production of TeV and PeV neutrinos by protons accelerated by the plasma shock wave of GRBs was discussed \cite{Meszaros 2001, Waxman 2007} and such high-energy neutrinos were searched for by AMANDA~\cite{Achterberg 2007}, ANITA~\cite{Vieregg 2011}, ANTARES~\cite{ANTARES1 2013, ANTARES2 2013}, Baikal~\cite{Avrorin 2009}, IceCube ~\cite{Abbasi 2010, Aartsen 2015}, and SuperKamiokande~\cite{Fukuda 2002}, but no signal was found. The searches for GRB neutrinos in the MeV energy range have been performed by four experiments: SuperKamiokande \cite{Fukuda 2002}, SNO \cite{Aharmim 2014}, KamLAND \cite{Asakura 2015}, and BUST~\cite{BAKSAN2015}. The SuperKamiokande collaboration searched for electron and muon neutrinos and antineutrinos in the energy range of 7--80\,MeV. The SNO collaboration searched for electron neutrinos, electron antineutrinos, and for (anti)neutrinos of non-electron flavors in the range of 5--13\,MeV. The KamLAND collaboration set upper limits on electron antineutrino fluence associated with GRBs with known redshift, in the energy ranges of 7.5--100~MeV and (after the Japanese nuclear reactors were switched off in 2011) of 1.8--100\,MeV. The BUST (Baksan Underground Scintillation Telescope) was sensitive to electron neutrinos and antineutrinos in the energy interval of 20--100\,MeV. None of these four experiments found any correlation between GRBs and neutrino events in their detectors. In this paper, we present a search for possible correlations between GRBs and (anti)neutrino events for all neutrino flavours in the Borexino detector. \begin{figure}[t] \begin{center} \includegraphics[width = 8 cm] {Detector.pdf} \caption{Schematic view of the Borexino detector.} \label{fig:Borexino} \end{center} \end{figure}
\label{sec:concl} We have performed a search for time correlation between gamma ray bursts and events detected by Borexino and associated with neutrinos and antineutrinos reactions -- the inverse beta decay reaction on protons and neutrino-electron elastic scattering. We have also searched for correlations of GRBs with short bursts of events in Borexino. The analysis was performed with data of two semi-independent data acquisition systems: the primary DAQ, optimized for events up to a few MeV, and a Flash ADC system, designed for events above 1\,MeV. A set of 2350 GRBs observed between December 2007 and November 2015 was checked for correlations with data acquired with the primary DAQ system. A set of 1813 GRBs between December 2009 and November 2015 were also checked for correlations with the data from the FADC DAQ system. We found no statistically significant time correlations of GRBs with the events in the detector in a time windows of $\pm$5000\,s around GRBs for $\bar{\nu}_e$-like events and for bursts of events, and of $\pm$1000\,s for neutrinos of all species. The demonstrated sensitivity in the MeV range is close to the neutrino fluences predicted by some models of GRBs. Limits on the fluence of neutrinos of all flavours (and, separately, of electron antineutrinos) were set for neutrino energies 1.5--15\,MeV. These are the most stringent bounds for GRB correlated fluence of neutrinos of all species below 7\,MeV, and on the GRB correlated fluence of $\bar{\nu}_e$'s in the range of 2--8\,MeV.
16
7
1607.05649
A search for neutrino and antineutrino events correlated with 2350 gamma-ray bursts (GRBs) is performed with Borexino data collected between December 2007 and November 2015. No statistically significant excess over background is observed. We look for electron antineutrinos (νbar<SUB>e</SUB>) that inverse beta decay on protons with energies from 1.8 MeV to 15 MeV and set the best limit on the neutrino fluence from GRBs below 8 MeV. The signals from neutrinos and antineutrinos from GRBs that scatter on electrons are also searched for, a detection channel made possible by the particularly radio-pure scintillator of Borexino. We obtain currently the best limits on the neutrino fluence of all flavors and species below 7 MeV. Finally, time correlations between GRBs and bursts of events are investigated. Our analysis combines two semi-independent data acquisition systems for the first time: the primary Borexino readout optimized for solar neutrino physics up to a few MeV, and a fast waveform digitizer system tuned for events above 1 MeV.
false
[ "Borexino data", "Borexino", "time correlations", "solar neutrino physics", "GRBs", "events", "November", "MeV", "bursts", "December", "the primary Borexino readout", "neutrinos", "two semi-independent data acquisition systems", "neutrino and antineutrino events", "a fast waveform digitizer system", "electron antineutrinos", "beta decay", "first", "a few MeV", "energies" ]
5.797765
-0.719511
13
12472413
[ "Yıldız, M.", "Çelik Orhan, Z.", "Kayhan, C." ]
2016MNRAS.462.1577Y
[ "Fundamental properties of Kepler and CoRoT targets - III. Tuning scaling relations using the first adiabatic exponent" ]
25
[ "Department of Astronomy, Space Sciences, Science Faculty, Ege University, 35100, Bornova, İzmir, Turkey", "Department of Astronomy, Space Sciences, Science Faculty, Ege University, 35100, Bornova, İzmir, Turkey", "Department of Astronomy, Space Sciences, Science Faculty, Ege University, 35100, Bornova, İzmir, Turkey" ]
[ "2017ApJ...843...11V", "2017MNRAS.470L..25Y", "2018ApJS..236...42Y", "2018ApJS..239...16F", "2019AJ....157..245H", "2019ApJ...879...33V", "2019ApJ...885...31C", "2019MNRAS.487.1335K", "2019MNRAS.489.1753Y", "2019MNRAS.490.1509K", "2020FrASS...7....3H", "2020NatAs...4..382C", "2021MNRAS.501.3162L", "2021MNRAS.503.4529C", "2021MNRAS.504.2273Y", "2021MNRAS.506.4413C", "2021NewA...8401522K", "2021RAA....21..226J", "2022ApJ...926..191Z", "2023A&A...669A..67H", "2023MNRAS.518.5552Y", "2023MNRAS.526.1799C", "2023Natur.618..917H", "2024ApJ...962...20W", "2024MNRAS.528.6881Y" ]
[ "astronomy" ]
7
[ "stars: evolution", "stars: fundamental parameters", "stars: interiors", "stars: late-type", "stars: oscillations", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1093/mnras/stw1709", "10.48550/arXiv.1607.03768" ]
1607
1607.03768.txt
%1\numino % %2.$\numino$ % %Stars are good. The best are solar-like osc. stars. Their advantage and usefullness in our understanding of everything. %--> kepler stars. %Physical processes are occuring deep inside stars in many different ways and kinds. {Many different types of physical processes occur deep inside stars.} The standard stellar models (SSMs), however, are constructed by taking into account only {the most basic} of the most essential physical principles (such as structure equations, matter-matter and matter-radiation interactions). Therefore, it is {inevitable} that our SSMs {will be} inadequate in some respects. {Progress} depends on {finding} discrepancies between SSMs and stars. %Therefore, it is expected %that the standart models constructed %by comprising the most essential rules such as stellar structure equations, energy production and transfer courses etc., %(nd our current knowledge on EOS, opacity and nuclear reaction rates) %are inadequate in some respects. Very detailed analysis is required for improvement in stellar physics, and also for discovering new processes occuring inside stars, {and only} very precise constraints can lead to {the discovery of such processes. Helioseismology and asteroseismology of solar-like oscillating stars are able to provide such constraints} %Such constraints are available from helioseismology and asteroseismology of solar-like oscillating stars (see e.g. Kosovichev 2011, Chaplin \& Miglio 2013, Christensen-Dalsgaard 2002, 2016). { Oscillation frequencies { ($\nu$)} of such stars are now available from the $Kepler$ (Borucki et al. 2010) and $CoRoT$ (Baglin et al. 2006) space missions, and from ground-based observations (Bedding et al. 2010, Bedding et al. 2007 and Bazot et al. 2012). According to scaling relations, stellar mass ($M$) and radius ($R$) can be found from frequency of maximum amplitude ($\nu_{\rm max}$), the large separation between the oscillation frequencies ($\Delta \nu$) and effective temperature (\teff). $\Delta \nu$ is not constant and {therefore} its mean value ($\braket{\Dnu}$) is used in these computations. Furthermore, it has an oscillatory component. {We have shown in two recent papers (\yildiz $ $ \etal 2014, hereafter Paper I; \yildiz, \c{C}elik Orhan \& Kayhan 2015, hereafter Paper II) %In two recent papers of us (\yildiz $ $ \etal 2014, hereafter Paper I; \yildiz, \c{C}elik Orhan \& Kayhan 2015, hereafter Paper II), we have shown that there is very strong diagnostic potential in the new reference frequencies ($\nu_{\rm min1}$ and $\nu_{\rm min2}$) at which \Dnu$ $ is minimum. } } %In the present study, we analyse observed oscillation frequencies inferred from $Kepler$ and $CoRoT$ photometry, and apply the methods %developed in Paper I and Paper II. In turn, we develop %similar methods by constructing interior models with the {\small MESA} code %(Paxton et al. 2011, 2013). %By this way, we test if the relations presented in Paper I and II are model dependent. {Paper I involved the investigations of models constructed using the {\small ANK\.I } code for the solar-like oscillating stars with solar composition, revealing for the first time new relations between oscillation frequencies and fundamental stellar parameters.} %are investigated and %new relations between oscillation frequencies and fundamental stellar parameters are obtained for the first time. The discovery of new reference frequencies $\nu_{\rm min1}$ and $\nu_{\rm min2}$ made these relations available. In order to derive general relations, the effects of metallicity ($Z$) and helium abundance ($Y$) should be clarified. In Paper II, we {attempted} to generalize the relations for { arbitrary $M$, $R$, $Z$ and $Y$ values}. %, and %develop new methods for determination of $Z$ and $Y$ from oscillatory components of the frequencies. %The effects of these abundances on the relations for stellar parameters in terms of asteroseismic %quantities are derived. % In the present study, we analyse observed oscillation frequencies, {confirm} the relations between the reference frequencies, and apply the new methods to the $Kepler$ and $CoRoT$ target stars. {Their \teff, $R$ and $M$ were found using} asteroseismic parameters. Such an application is also very important for testing the scaling relations. %These parameters seem to form a complete set (as similar to %base vectors in wave function space) in the stellar parameters space. %\textcolor{red}{scaling relation literature summary}:\\ In derivation of the scaling relations used to compute { $M$ and $R$ in terms of \numax, $\braket{\Dnu}$ and $\teff$, } it is assumed that the first adiabatic exponent ($\Gamma_{\rm \negthinspace 1s}$) and mean molecular weight ($\mu$) at stellar surface are constant (Brown et al. 1991; Kjeldsen \& Bedding 1995). We test {whether} these quantities are constant {for our purposes and, if not, recommend that the scaling relations should be verified (see Section 3).} The relation { between $\braket{\Dnu}$ and } mean density ($\braket{\rho}$) {has been} widely discussed in recent papers. White et al. (2011) state that the $\braket{\Dnu}$-$\braket{\rho}$ relation depends { on \teff$ $ and} suggest a fitting formula for correction in order to find $\braket{\rho}$ from the observed values of $\braket{\Dnu}$. In addition to the $\braket{\Dnu}$-$\braket{\rho}$ relation, Belkacem et al. (2013) also discuss the relation between \numax$ $ and acoustic cut-off frequency ($\nu_{\rm ac}$), {claiming} that departure from the observed relation arises from the complexity of non-adiabatic processes. Garc{\'{i}}a Hern{\'{a}}ndez et al. (2015), however, derive an observational scaling ($\braket{\Dnu}$-$\braket{\rho}$) relation for the $\delta$ Scuti components in eclipsing binaries. Recently, Sharma et al. (2016) {aimed} to generalize problems pertaining to the scaling relations for \kepler$ $ red giants, without considering $\Gamma_{\rm \negthinspace 1s}$ as variable. Our strategy is first to {identify whether} $constants$ are constant, {before attempting to find} parameters for $\braket{\Dnu}$-$\braket{\rho}$ and \numax-$\nu_{\rm ac}$ relations. %?A very important discrepancy exists between model and observed radii of late-type component stars in detached eclipsing %binaries. Details of literature. Non-ideal effects. K enhancement. If there is any evidence for this discrepancy. For the assessment of the results on fundamental properties of stars, it is {crucial to derive the observed values of these parameters by alternative direct methods, for which, the roles of Sun and { Procyon A} are of key importance.} %For now, there are two such stars. One is the Sun. The other star is Procyon. From astrometric observations of Procyon by $Hubble$ $Space$ $Telescope$, {the} mass of its primary component is determined very precisely, $1.478 \pm 0.012$ \MS (Bond et al. 2015). This data {enables the testing of} new scaling relations (see Section 4). In addition to interferometrically observed solar-like oscillating stars (Huber et al. 2011b, Baines et al. 2014), component stars in eclipsing binaries are benchmark for asteroseismic studies (Gaulme et al. 2013, Rawls et al. 2016), despite the {complicating factor of} tidally induced oscillations. This paper is organized as follows. {Section 2 is the presentation of the basic asteroseismic and non-asteroseismic properties of the target stars compiled from literature.} In this section, we also compare $\nu_{\rm min1}$ and $\nu_{\rm min2}$ of the models with their observational counterparts. Section 3 is devoted to {\small MESA} { (Paxton et al. 2011)} models and role of $\Gamma_{\rm \negthinspace 1s}$ in new scaling relations. In this section, we also develop new expressions for effective temperature by using oscillation frequencies. In Section 4, {the results based on asteroseismic methods are presented and compared} with results obtained by conventional methods. Finally, in Section 5, {conclusions are drawn}.
%Formation and evolutionary scenario of scales from planetary systems to galactic structure require %precise information on fundamental properties of stars. { Asteroseismology of solar-like oscillating stars provides information about fundamental properties of these stars} through the scaling relations. In these relations it is assumed that the first adiabatic exponent $\Gamma_{\rm \negthinspace 1s}$ is constant. However, analysis of our models constructed by the {\small MESA} code shows that $\Gamma_{\rm \negthinspace 1s}$ significantly changes through the surfaces of solar-like oscillating stars, depending on effective temperature. Furthermore, the ratio of the { mean} large separation $\braket{\Dnu}$ to $\braket{\rho}^{0.5}$, which is customarily assumed to be constant, is a linear function of $\Gamma_{\rm \negthinspace 1s}$. Thus, it seems that $\Gamma_{\rm \negthinspace 1s}$ is the {factor in scaling relation that has previously been overlooked.} %missing piece of scaling relation puzzle. In contrast to the literature, we do not consider the ratios of $\nu_{\rm ac}/\numax $ and $\nu_{\rm ac}/(g/T_{\rm eff}^{0.5}) $ as { constant}. We show that these ratios can also be taken as functions of $\Gamma_{\rm \negthinspace 1s}$. Then, we revise the scaling relations and obtain new relations for stellar mass and radius { (equations 9 and 10)}, and then compute mass and radius of {\it Kepler} and {\it CoRoT} targets (89 stars + the Sun) using their asteroseismic properties. {A difference of up to 10 per cent exists between the masses from new and old scaling relations.} %The difference between masses from new and old scaling relations is up to 10 per cent. However, the great part of the uncertainties in mass and radius found from scaling relations comes from uncertainty in $\numax$. When {mass and radius obtained from scaling relations are compared with those available in the literature, significant differences are seen}. The fractional differences between the masses are up to $\Delta M/M = 0.25$ and $\Delta R/R = 0.10$. {Particularly noteworthy, however, is the linear relation between $\Delta M/M$ and $\Delta R/R$: $\Delta M/M = 3 \Delta R/R$ { (see Fig. 14)}, highlighting that the fitted parameter is indeed mean density when oscillation frequencies of interior models are fitted to the observed oscillation frequencies.} %it is very interesting that there is a linear relation between $\Delta M/M$ and $\Delta R/R$: $\Delta M/M = 3 \Delta R/R$ { (see Fig. 14)}. %This means that the fitted parameter is indeed mean density %when one fits oscillation frequencies of interior models to the observed oscillation frequencies. We also {computed effective temperatures of these stars using} purely asteroseismic methods. In our previous paper (Paper I), {the effective temperature of models was shown to be a function of $\Delta n_{\rm x1}$, which is approximately} the order difference between the { frequencies} of maximum amplitude and min1 in the $\Dnu - \nu$ graph. {Taking} into account the effect of $\Gamma_{\rm \negthinspace 1s}$, we derive new relations between $T_{\rm eff}$ and $\Delta n_{\rm x1}$ by using the {\small MESA} models. Similar expressions are obtained for $\Delta n_{\rm x0}$ and $\Delta n_{\rm x2}$. {This allows us in principle to use three different methods to compute $T_{\rm eff}$} in terms of the oscillation frequencies of the target stars. %Then, %we can in principle compute $T_{\rm eff}$ from three different methods in terms of the oscillation %frequencies of the target stars. i These effective temperatures ($T_{\rm sis0}$, $T_{\rm sis1}$ and $T_{\rm sis2}$; see { equations 15-17}) are in general in very good agreement within themselves and with $T_{\rm eff}$ from conventional methods. {A significant} difference appears between $T_{\rm sis1}$ and $T_{\rm eS}$ {for cool stars, but not for the Sun, for example.} {The value of these new approaches lies in the increased number of methods they allow for computing the effective temperature of a solar-like oscillating star. The six methods consist of three asteroseismic, one spectroscopic and two photometric methods.} %These new methods are very important. Thanks to them, %we can compute effective temperature of a solar-like oscillating star with 6 different methods (three asteroseismic, one spectroscopic and two %photometric methods). In principle, we can compute the fundamental stellar parameters by purely asterosesimic quantities, provided that the oscillation frequencies are precisely determined. The solar effective temperature is found as 5742, 5831 and 5840 K from frequencies of min0, min1 and min2, respectively. All of these values are very close to $T_{\rm eff \odot}$. Another {key} result we obtain is about { Procyon A}. Its mass is determined by using asterometric data from $Hubble$ $Space$ $Telescope$ as $1.478\pm0.012$ \MSbit. While the conventional scaling relation yields 1.63 \MSbit, the new scaling relation gives {a much more accurate figure,} 1.46 \MSbit. However, in some cases, we confirm that there are systematic differences between asteroseismic and non-asteroseismic effective temperatures. These differences can be {reduced or eliminated} by increasing or decreasing \numax. Such a {modification, ciritical} for obtaining more precise mass and radius from the scaling {relations, will} be the subject of our next paper. %The stars we study are the stars for which the oscillation frequencies are available in the literature. %Apporchoux et al. (2012) give the frequencies and $\numax$ of 31 stars. We determine $\nu_{\rm min1}$ and $\nu_{\rm min2}$ %of these stars and then compute the masses of these stars by using different versions of asteroseismic expressions %for the mass and other stellar properties. The results of our computations are given in Table 1. In the present study we have investigated the physical properties of CBs. The approximate initial masses of the secondary and primary components are computed, for the first time ever, by using the basic parameters of 100 W UMa type CBs (51 A- and 49 W-subtype). The initial mass of the secondary is computed from its luminosity excess by a method based on the expression derived from stellar models with mass-loss. We have also deduced that only one-third of the mass lost by the secondary components is transferred to the primary components. The remaining part leaves the systems. The range of initial masses of the secondaries of CBs is $1.3-2.6 \mathrm{M}_{\odot}$. The upper part of this range is occupied by the A-subtype CBs and the W-subtype CBs are in the lower part. The transition occurs at about $1.8\pm 0.1\mathrm{M}_\odot$. %This result leads us to confirm that %detached binary systems with an early-type secondary (the component that was {\it more} massive before %%the mass transfer process) of mass $M_{\rm 2i} <$ 2.6 M$_\odot$ become A-subtype CBs, while the W-subtype %CBs are mostly formed from the detached binaries with late-type components. %\par\noindent ****** I didn't understand your last sentence of this paragraph, but it seems to me to be %only saying in a different form what you have already said. ********* The main assumption in our method is that the mass transfer starts near or after the TAMS phase of the massive component (progenitor of the secondary component). This assumption seems fairly good for the A-subtypes. For the W-subtype, however, a correction is needed. From $\delta M$ and the reciprocal mass-ratio diagram (Fig. 7), we develop a method and obtain corrected initial masses for the case of small $\delta M$. The corrected initial masses of the secondaries range from $1.3-1.9 \mathrm{M}_{\odot}$. One of the important outcomes of the present study is that there is also a lower limit for the initial mass of the primary components. The reason for this result should be studied in the context of angular momentum evolution versus nuclear evolution. For the binary systems with an initial mass higher than $1.8 \mathrm{M}_{\odot}$, although the angular momentum is lost by the less massive component the nuclear evolution is principally the mechanism responsible for the Roche lobe filling process. Therefore, semi-detached systems with $M_{\rm 2i}=1.8-2.61 \mathrm{M}_{\odot}$ must yield A-subtype W UMa CBs. The binary systems with $M_{\rm 2i}$ less than $1.8 \mathrm{M}_{\odot}$ evolve into the contact phase due to the rapid angular momentum evolution. The Roche lobes of their components contract so that they form the W-subtype contact or near-contact binaries despite the components not being much evolved in the MS phase. The primary components of W UMa type CBs have an initial mass range from $0.2$ to $1.5 \mathrm{M}_{\odot}$. %\par\noindent ******* Do you mean the initial version of the current primary, i.e the component that was %initially the secondary system? I find this terminology rather confusing, and I suspect others will. %Although it's a bit late, I would favour A and B as the names of the {\it currently} more and less %massive components, and $*1$ and $*2$, pronounced `star 1' and `star 2', as the names of the {\it initially} %more and less massive components. ************ % %\par\noindent A star with mass in this range loses its spin angular momentum relatively rapidly. If it has a tidally interacting companion, orbital angular momentum is lost by the system. The precursors of W UMa type CBs are the systems in which nuclear evolution of the massive components is in competition with the angular momentum loss from the less massive components. In the precursors of the W-subtype W UMa CBs, both components are effective in the angular momentum loss process. The method we develop for computation of the initial masses of the secondary components is highly suitable for a complete error analysis. From the derivatives of the function derived from the fitting curve for $\Delta M$ as a function of $L_2$ and $M_2$ (equation 8), we successfully compute $\Delta M_{\rm 2i}$ in terms of the observed uncertainties $\Delta L_2$ and $\Delta M_2$ (equation 11). We also consider cluster member W UMa type CBs. Although the number of CBs with accurate dimensions available in the literature is very small (4 systems), the results support our findings on initial masses and time-scales for the A- and W-subtype CBs.
16
7
1607.03768
So-called scaling relations based on oscillation frequencies have the potential to reveal the mass and radius of solar-like oscillating stars. In the derivation of these relations, it is assumed that the first adiabatic exponent at the surface (Γ_{1s}) of such stars is constant. However, by constructing interior models for the mass range 0.8-1.6 M<SUB>⊙</SUB>, we show that Γ _{1s} is not constant at stellar surfaces for the effective temperature range with which we deal. Furthermore, the well-known relation between large separation and mean density also depends on Γ _{1s}. Such knowledge is the basis for our aim of modifying the scaling relations. There are significant differences between masses and radii found from modified and conventional scaling relations. However, a comparison of predictions of these relations with the non-asteroseismic observations of Procyon A reveals that new scaling relations are effective in determining the mass and radius of stars. In the present study, solar-like oscillation frequencies of 89 target stars (mostly Kepler and CoRoT) were analysed. As well as two new reference frequencies (ν<SUB>min1</SUB> and ν<SUB>min2</SUB>) found in the spacing of solar-like oscillation frequencies of stellar interior models, we also take into account ν<SUB>min0</SUB>. In addition to the frequency of maximum amplitude, these frequencies have a very strong diagnostic potential in the determination of fundamental properties. The present study applies the derived relations from the models to the solar-like oscillating stars, and computes their effective temperatures using purely asteroseismic methods. There are in general very close agreements between effective temperatures from asteroseismic and non-asteroseismic (spectral and photometric) methods. For the Sun and Procyon A, for example, the agreement is almost total.
false
[ "effective temperatures", "new scaling relations", "non-asteroseismic", "stellar surfaces", "such stars", "oscillation frequencies", "stellar interior models", "stars", "asteroseismic", "fundamental properties", "solar-like oscillating stars", "_", "solar-like oscillation frequencies", "interior models", "the effective temperature range", "Procyon A", "Γ", "modified and conventional scaling relations", "the non-asteroseismic observations", "radii" ]
7.559485
11.70316
81
12513984
[ "Arkani-Hamed, Nima", "Cohen, Timothy", "D'Agnolo, Raffaele Tito", "Hook, Anson", "Kim, Hyung Do", "Pinner, David" ]
2016PhRvL.117y1801A
[ "Solving the Hierarchy Problem at Reheating with a Large Number of Degrees of Freedom" ]
134
[ "School of Natural Sciences, Institute for Advanced Study, Princeton, New Jersey 08540, USA", "Institute of Theoretical Science, University of Oregon, Eugene, Oregon 97403, USA", "School of Natural Sciences, Institute for Advanced Study, Princeton, New Jersey 08540, USA", "Stanford Institute for Theoretical Physics, Stanford University, Stanford, California 94305, USA", "Department of Physics and Astronomy and Center for Theoretical Physics, Seoul National University, Seoul 151-747, Korea", "Princeton Center for Theoretical Science, Princeton University, Princeton, New Jersey 08544, USA" ]
[ "2016JHEP...12..101H", "2016arXiv161002743A", "2017JCAP...11..007B", "2017JHEP...05..038C", "2017JHEP...06..126K", "2017JHEP...07..023C", "2017JHEP...11..125H", "2017JHEP...12..139B", "2017MPLA...3250049A", "2017PhRvD..96b3501A", "2017PhRvD..96e5013C", "2017arXiv170104338H", "2017arXiv170305733D", "2017arXiv170501472T", "2017arXiv170509698T", "2017arXiv170605008H", "2017arXiv170900766D", "2017arXiv171007663G", "2017arXiv171205926H", "2018JCAP...01..022B", "2018JCAP...06..044C", "2018JHEP...07..033F", "2018JHEP...10..014G", "2018JHEP...10..151F", "2018JHEP...12..094D", "2018PDU....19....1C", "2018PDU....21...21K", "2018PhDT.......146W", "2018PhLB..785..585Y", "2018PhRvD..97e5030H", "2018PhRvD..98d3503M", "2018PhRvL.120z1802H", "2018arXiv180509345H", "2018arXiv181200637H", "2019BAAS...51c.159G", "2019EPJC...79..862L", "2019JHEP...01..081C", "2019JHEP...04..027G", "2019JHEP...05..072L", "2019JHEP...06..112L", "2019JHEP...08..151A", "2019JHEP...10..060S", "2019JHEP...10..199G", "2019JPhCS1147a2083G", "2019PhLB..788..288H", "2019PhRvD..99a5005H", "2019PhRvD..99b3529W", "2019PhRvD..99c1702C", "2019PhRvD..99d3513D", "2019PhRvD..99e5047B", "2019PhRvD..99i5026W", "2019PhRvD..99j3526K", "2019PhRvD.100c5006C", "2019PhRvD.100i5001G", "2019PhRvL.122s1802G", "2019arXiv190303622C", "2019arXiv190704473A", "2020EPJC...80..124L", "2020EPJC...80.1188C", "2020IJMPA..3550182N", "2020JCAP...12..025H", "2020PhRvD.101h3533D", "2020PhRvD.101i5016A", "2020PhRvD.101k5002S", "2020PhRvD.101k5007M", "2020PhRvD.102b6012C", "2020PhRvD.102i5004G", "2020PhRvD.102j3512K", "2020arXiv200500292N", "2020arXiv200810625H", "2020arXiv200907817B", "2020arXiv200911870K", "2020arXiv201012459B", "2021EPJC...81..388B", "2021JCAP...05..065K", "2021JHEP...10..093G", "2021PhRvD.103c5005E", "2021PhRvD.103e5014D", "2021PhRvL.126i1801C", "2021arXiv210604591T", "2021arXiv210913249T", "2021arXiv211101096G", "2021lhcp.confE..19G", "2022ApJ...934..122R", "2022EPJWC.27408011O", "2022JCAP...02..019G", "2022JCAP...07..034B", "2022JCAP...10..046Y", "2022JHEP...02..023D", "2022JHEP...03..113B", "2022JHEP...03..114C", "2022JHEP...05..050B", "2022JHEP...07..126D", "2022JHEP...08..266B", "2022MNRAS.516.2038B", "2022Natur.607...41S", "2022PhLB..82736991B", "2022PhRvD.105c5008J", "2022PhRvL.128b1803D", "2022PhRvL.128t1301C", "2022PhRvR...4b2048J", "2022RPPh...85h4201A", "2022ScPP...12..150B", "2022arXiv220305010E", "2022arXiv220305531B", "2022arXiv220306680A", "2022arXiv220307354B", "2022arXiv220307622A", "2022arXiv220307943D", "2022arXiv220508553C", "2022arXiv221003075F", "2022arXiv221208685G", "2023ARNPS..73...23H", "2023EPJC...83..825C", "2023JHEP...01..063B", "2023JHEP...02..142E", "2023JHEP...04..082C", "2023JHEP...06..107C", "2023JPhG...50c0501F", "2023PDU....4201333G", "2023PhRvD.107b3517G", "2023PhRvD.107k5011B", "2023PhRvD.107l3517B", "2023PhRvD.108k5019H", "2024ApJ...965...42A", "2024EL....14639002W", "2024JCAP...01..064F", "2024JHEP...01..148B", "2024JHEP...02..021C", "2024JHEP...06..052B", "2024PhRvD.109g5003E", "2024PhRvD.109i5016E", "2024arXiv240305619G", "2024arXiv240500836C" ]
[ "astronomy", "physics" ]
16
[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]
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[ "10.1103/PhysRevLett.117.251801", "10.48550/arXiv.1607.06821" ]
1607
1607.06821_arXiv.txt
This letter describes a new mechanism, dubbed ``$N$naturalness,'' which solves the hierarchy problem. It predicts no new particles at the LHC, but does yield a variety of experimental signatures for the next generation of CMB and large scale structure experiments~\cite{Dodelson:2013pln, Allison:2015qca}. Well-motivated supersymmetric incarnations of this model predict superpartners beneath the scale $m_W \times M_{\text{pl}}/M_{\text{GUT}} \sim 10$~TeV, accessible to a future 100 TeV collider~\cite{Arkani-Hamed:2015vfh, Golling:2016gvc}. The first step is to introduce $N$ sectors which are mutually non-interacting. The detailed particle content of these sectors is unimportant, with the exception that the Standard Model (SM) should not be atypical; many sectors should contain scalars, chiral fermions, unbroken gauge groups, etc. For simplicity, we imagine that they are exact copies of the SM, with the same gauge and Yukawa structure. It is crucial that the Higgs mass parameters are allowed to take values distributed between $-\Lambda_H^2$ and $\Lambda_H^2$, where $\Lambda_H$ is the (common) scale that cuts off the quadratic divergences. Then for a wide range of distributions, the generic expectation is that some sectors are accidentally tuned at the $1/N$ level, $\left|m_H^2\right|_{\text{min}} \sim \Lambda_H^2 / N$. We identify the sector with the smallest non-zero Higgs vacuum expectation value (vev), $\vev{H} = v$, as ``our" SM. This picture is illustrated schematically in Fig.~\ref{fig:Sketch}. \begin{figure}[!t] \centering \includegraphics[width=0.48\textwidth]{Figures/N_Copies_Sketch.pdf} \caption{A sketch of the $N$naturalness setup. The sectors have been ordered so that they range from $m_H^2 \sim \Lambda_H^2$ to $-\Lambda_H^2$. The sector with the smallest vacuum expectation value contains our copy of the SM.} \label{fig:Sketch} \end{figure} In order for small values of $m_H^2$ to be populated, the distribution of the mass parameters must pass through zero. For concreteness, we take a simple uniform distribution of mass squared parameters, indexed by an integer label $i$ such that \be \left(m_H^2\right)_i = -\frac{\Lambda_H^2}{N} \big(2\, i+ r \big),\quad\quad -\frac{N}{2} \leq i \leq \frac{N}{2}, \label{eq:mHsqiScaling} \ee where $i=0 = \text{``us"}$ is the lightest sector with a non-zero vev: $\left(m_H^2\right)_\text{us} = -r\times\Lambda_H^2/N \simeq -(88 \text{ GeV})^2$ is the Higgs mass parameter inferred from observations. The parameter $r$ can be seen as a proxy for fine-tuning,\footnote{There are a variety of other ways one might choose to implement a measure of fine-tuning in this model. For example, one could assume the distribution of Higgs mass squared parameters is random with some (arbitrary) prior, and then ask statistical questions regarding how often the resulting theory is compatible with observations.} since it provides a way to explore how well the naive relation between the cutoff and the mass scale of our sector works in a detailed analysis. Specifically, $r = 1$ corresponds to uniform spacing, while $r<1$ models to an accidentally larger splitting between our sector and the next one. A simple physical picture for this setup is that the new sectors are localized to branes which are displaced from one another in an extra dimension. In this scenario, the lack of direct coupling is clear, and the variation of the mass parameters can be explained geometrically: the $m_H^2$ parameters may be controlled by the profile of a quasi-localized field shining into the bulk. As a consequence of the existence of a large number of degrees of freedom, the hierarchy between $\Lambda_H$ and the scale $\Lambda_G$ where gravity becomes strongly coupled is reduced. The renormalization of the Newton constant implies $\Lambda_G^2 \sim M_{\text{pl}}^2 / N$. If perturbative gauge coupling unification is to be preserved $\Lambda_G \gtrsim M_{\text{GUT}}$, implying that $N \lesssim 10^4$. This gives a cutoff no greater than $\Lambda_H \sim 10$~TeV, thus predicting a little hierarchy that mirrors the GUT-Planck splitting in the UV. At the scale $\Lambda_H$, new dynamics (\emph{e.g.}, SUSY) must appear to keep the Higgs from experiencing sensitivity to even higher scales. Alternatively, the full hierarchy problem can be solved with $N \sim 10^{16}$, so that $\Lambda_H \sim \Lambda_G \sim 10^{10}$~GeV. Note that this number of copies, while sufficient, is unnecessary for a complete solution. There may be two classes of new degrees of freedom: the $N$ copies that participate directly in the $N$naturalness picture, and another completely sterile set of degrees of freedom that still impact the renormalization of $\Lambda_G$. So far we have described a theory with a $S_N$ permutation symmetry, broken softly by the $m_H^2$ parameters, such that each of the sectors is SM-like. Sectors for which $m_H^2 < 0$ are similar to our own, with the exception that particle masses scale with the Higgs vev, $v_i \sim v\,\sqrt{i}$. In addition, once $i \gtrsim 10^{8}$ the quarks are all heavier than their respective QCD scales. Those sectors do not exhibit chiral symmetry breaking, nor do they contain baryons. Sectors with $m_H^2>0$ are dramatically different from ours. In these sectors, electroweak symmetry is broken at low scales due to the QCD condensate $\Lambda_\text{QCD}$. Fermion masses are generated by the four-fermion interactions that are induced by integrating out the complete $SU(2)$ Higgs multiplet. Thus, $m_f \sim y_f\,y_t\,\Lambda_{\text{QCD}}^3 / \left(m_{H}^2\right)_i \lesssim 100$~eV, where $y_t$ is the top Yukawa coupling. All fermionic and gauge degrees of freedom are extremely light relative to the ones in our sector. With so many additional degrees of freedom, the naive cosmological history is dramatically excluded. In particular, if all sectors have comparable temperatures in the early Universe, then one expects $\Delta N_{\text{eff}} \sim N$ (see Eq.~(\ref{eq:dNeff})). Thus, the hierarchy problem gets transmuted into the question of how to predominantly reheat only those sectors with a tuned Higgs mass. To accomplish this, we need to introduce a last ingredient into the story, the ``reheaton" field, so named because it is responsible for reheating the Universe via its decays. We call this field $S^c$ for models where the reheaton is a fermion, and $\phi$ if the reheaton is a scalar. The cosmological history of the model begins in a post-inflationary phase where the energy density of the Universe is dominated by the reheaton. As stated multiple times we can not be unique, therefore we assume that the reheaton couples universally to all sectors. Note that the scalars must be near their true minimum when reheating occurs. This can be accomplished by having either low scale inflation, or else a coupling of the Higgses to the Ricci scalar. In the next section, we present a set of models in which the reheaton dynamically selects and populates only the lightest sectors, despite preserving the aforementioned softly broken $S_N$ symmetry. Sec.~\ref{sec:Signals} then provides constraints on these models, and Sec.~\ref{sec:Discussion} contains our conclusions and highlights potential signals.
\label{sec:Discussion} In this paper we have proposed a new solution to the hierarchy problem. The need for a huge integer $N$ is obviously the least appealing feature of our setup. It is perhaps not entirely unreasonable to have the mild $N \sim 10^4$ compatible with the existence of a supersymmetric GUT scale, but this seems outlandish in the $N \sim 10^{16}$ limit. At the moment it is difficult to see how such a large integer can be explained dynamically, in the same way as we usually explain hierarchies by, \emph{e.g.} dimensional transmutation. On the other hand, this is simply another large set of degrees of freedom, and we do not deeply understand where the even vaster number of degrees of freedom in a macroscopic expanding universe comes from, so perhaps the large $N$ may eventually find a different sort of natural explanation. The theoretical consistency of the proposal also makes a number of demands on the UV theory, such as the absence of sizable cross-couplings between the sectors, which may be technically natural but may again strain credulity. However, we find it fascinating that huge values of $N$ are experimentally viable. This is highly non-trivial, and indeed in the simplest models we did find significant constraints on $N$. While we have examined all the zeroth-order phenomenological constraints we know of, it is important to continue to look for constraints on (and signals of!)\! the scenarios with high values of $N \big(\gg 10^4\big)$. It is also interesting to compare $N$naturalness with other approaches. It bears a superficial resemblance to large extra dimensions, which add $10^{32}$ degrees of freedom in the form of KK gravitons, as well as the scenario of Dvali~\cite{Dvali:2007hz} which invokes $10^{32}$ copies of the SM. In each of these cases, $M_{\text{pl}}$ is renormalized down to the TeV~scale. Of course this predicts (as yet unseen) new particles accessible to the LHC~\cite{Dvali:2009ne}. By contrast, $N$naturalness solves the hierarchy problem with cosmological dynamics; the weak scale is parametrically removed from the cutoff, and so it does not demand new physics to be accessible at colliders. $N$naturalness has some features in common with low-energy SUSY as well. Both models invoke a softly broken symmetry: SUSY is broken by soft terms, and the $S_N$ symmetry is broken by varying Higgs masses. Also in both cases, the most obvious implementations of the idea are experimentally excluded. If SUSY is directly broken in the MSSM sector, we have the famous difficulties with charge and color breaking; in the case of $N$naturalness, direct reheating of all $N$ sectors is grossly excluded by $N_\text{eff}$. Thus in both cases we need to have ``mediators.'' SUSY must be dominantly broken in another sector and have its effects mediated to the MSSM. Similarly, reheating must be dominantly communicated to the reheaton, which subsequently dumps its energy density into the other sectors. Finally, both models have additional scales that are not, on the face of it, tied to the physics responsible for naturalness. In SUSY there is a ``$\mu$ problem'' in that the vector-like Higgsino mass must be comparable to the soft scalar masses, while in $N$naturalness the reheaton mass must be close to the bottom of the spectrum of Higgs masses. While in both cases there are simple pictures for how this can come about, these coincidences do not emerge automatically. Moving beyond purely field theoretic mechanisms, there is the recent proposal of the relaxion~\cite{Graham:2015cka}, which invokes an extremely long period of inflation coupled with axionic dynamics to relax to a low weak scale. While both the relaxion and $N$naturalness mechanisms are cosmological, the physical mechanism of the relaxation, associated with the huge number of $e$-foldings of inflation, is {\it in principle} unobservable given our current accelerating Universe, much like the vast regions of the multiverse outside our cosmological horizon are imperceptible. By contrast, the cosmological dynamics associated with reheaton decay in $N$naturalness are sharply imprinted on the particle number abundance in all the sectors. They are not only in principle observable but, as we have stressed (at least for a small number of sectors ``close'' to ours), are detectable in practice within our Universe. It is also interesting to contrast $N$naturalness with the picture of an eternally inflating multiverse, with environmental selection explaining the smallness of the cosmological constant, as well as potentially at least part of the hierarchy problem. This picture is, after all, the first cosmological approach to fine-tuning puzzles. While it is very far from well-understood and has yet to make internal theoretical sense, it is the only cartoon we have for understanding the cosmological constant problem and does not involve any model-building gymnastics. Furthermore, fine-tuning for the Higgs mass also has a plausible environmental explanation. Especially in the context of minimal split SUSY~\cite{ArkaniHamed:2004fb}, these ideas give us a picture which simultaneously accounts for the apparent fine-tuning of the cosmological constant and the Higgs mass, while maintaining the striking quantitative successes of natural SUSY theories in the form of gauge coupling unification and dark matter. Nonetheless, it is important to continue to look for alternatives, minimally as a foil to the landscape paradigm. $N$naturalness is a concrete example of an entirely different cosmological approach to tuning puzzles, and in particular relies on the existence of only a single vacuum. We note that there is no obstacle to augmenting $N$naturalness with an anthropic solution to the cosmological constant problem. The presence of extra sectors exponentially increases the number of available vacua. For example we could add to the SM a sector with $m$ vacua and end up with $m^N$. Already $N\simeq 10^4$ with two vacua per sector is more than enough to scan the cosmological constant without relying on string theory landscapes. When solving the entire hierarchy problem with $N \simeq 10^{16}$, the vacua utilized to scan the cosmological constant can even be the two minima of the Higgs potential; this requires a high cutoff so that the second minimum is below $\Lambda_H$ and the difference in the potential energy of the two minima is $\mathcal{O}\big(\Lambda_G\big)$. To conclude, we would like to comment on the nature of the signals that we have discussed in this paper. For concreteness, three models that make $N$naturalness cosmologically viable were presented. However, it is easy to imagine a broader class of theories that realizes the same mechanism. We can relax the assumption that the Higgs masses are uniformly spaced (or even pulled from a uniformly distribution) or that all the new sectors are exact copies of the SM. It is also possible to construct different models of reheating, with new physics near the weak scale to modify the UV behavior of the theory. Nonetheless our sector can not be special in any way. There will always be a large number of other sectors with massless particles and with matter and gauge contents similar to ours, leading to the following signatures: \mbox{} \\ \begin{itemize}[leftmargin=12pt] \item We expect extra radiation to be observable at future CMB experiments. \item The neutrinos in the closest $m_H^2<0$ sectors are slightly heavier and slightly less abundant than ours. This implies $\mathcal{O}(1)$ changes in neutrino cosmology, which will start to be probed at this level in the next generation of CMB experiments~\cite{Allison:2015qca}. \item If the strong CP problem is solved by an axion, its mass will be much larger than the standard prediction. \item If $N\lesssim 10^4$ as motivated by grand unification, supersymmetry or new natural dynamics should appear beneath $10$~TeV. \end{itemize} \mbox{} \\ The natural parameter space is being probed now, and soon we may know if the $N$naturalness paradigm explains how the hierarchy problem has been solved by nature. \vspace{.3in}
16
7
1607.06821
We present a new solution to the electroweak hierarchy problem. We introduce N copies of the standard model with varying values of the Higgs mass parameter. This generically yields a sector whose weak scale is parametrically removed from the cutoff by a factor of 1 /√{N }. Ensuring that reheating deposits a majority of the total energy density into this lightest sector requires a modification of the standard cosmological history, providing a powerful probe of the mechanism. Current and near-future experiments can explore much of the natural parameter space. Furthermore, supersymmetric completions that preserve grand unification predict superpartners with mass below m<SUB>W</SUB>M<SUB>pl</SUB>/M<SUB>GUT</SUB>∼10 TeV .
false
[ "TeV", "mass", "the natural parameter space", "varying values", "the Higgs mass parameter", "grand unification", "m", "SUB>W</SUB>M<SUB>pl</SUB>/M<SUB>GUT</SUB>∼10 TeV", "Higgs", "superpartners", "the standard cosmological history", "the electroweak hierarchy problem", "a powerful probe", "the total energy density", "supersymmetric completions", "this lightest sector", "the mechanism", "deposits", "the standard model", "whose weak scale" ]
10.146914
-1.409307
-1
3481653
[ "Gheller, C.", "Vazza, F.", "Brüggen, M.", "Alpaslan, M.", "Holwerda, B. W.", "Hopkins, A. M.", "Liske, J." ]
2016MNRAS.462..448G
[ "Evolution of cosmic filaments and of their galaxy population from MHD cosmological simulations" ]
42
[ "ETHZ-CSCS, Via Trevano 131, CH-6900 Lugano, Switzerland", "Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112, D-21029 Hamburg, Germany; INAF-Istituto di Radio Astronomia, Via Gobetti 101, I-40129 Bologna, Italy", "Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112, D-21029 Hamburg, Germany", "NASA Ames Research Center, N232, Moffett Field, Mountain View, CA 94035, USA", "University of Leiden, Sterrenwacht Leiden, Niels Bohrweg 2, NL-2333 CA Leiden, The Netherlands", "Australian Astronomical Observatory PO Box 915, North Ryde NSW 1670, Australia", "Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112, D-21029 Hamburg, Germany" ]
[ "2017A&A...602L...6V", "2017CQGra..34w4001V", "2017MNRAS.467.4914V", "2017PhDT........11N", "2017PhRvL.119j1101M", "2018Galax...6..128L", "2018JHEAp..20....1G", "2018MNRAS.478.4336M", "2018MNRAS.479..776V", "2018MNRAS.479.3343M", "2018MNRAS.480.3749G", "2018PhRvD..98d3018G", "2018SSRv..214..122D", "2019A&A...627A...5V", "2019AJ....157..254L", "2019EPJC...79..921D", "2019MNRAS.486..981G", "2019MNRAS.486.3766M", "2019SSRv..215...16V", "2020A&A...637A..18B", "2020A&A...638A..75B", "2020A&A...641A.173G", "2020MNRAS.491.5447V", "2020MNRAS.491.5747M", "2020MNRAS.494.5603G", "2020MNRAS.495.4475M", "2021A&A...649A.117G", "2021A&A...652A..80L", "2021AJ....161..255E", "2021JCAP...02..048A", "2021MNRAS.502..351R", "2022A&A...661A..17B", "2022A&A...661A.115G", "2022A&A...662A..87O", "2022MNRAS.514.5429C", "2022MNRAS.516..231S", "2022Univ....8..253G", "2023A&A...675A..63E", "2023MNRAS.518..240O", "2023MNRAS.526..224W", "2024A&A...684A..63G", "2024arXiv240218455A" ]
[ "astronomy" ]
10
[ "methods: numerical", "intergalactic medium", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1093/mnras/stw1595", "10.48550/arXiv.1607.01406" ]
1607
1607.01406_arXiv.txt
\label{sec:intro} The large-scale structure of the Universe is organized into a web of filamentary matter connecting clusters and groups of galaxies and wall--like structures separating gigantic underdense regions (voids). A large fraction of the baryonic matter (around 50\%) is predicted to reside in such {\it cosmic web}, at densities $\sim 10-100$ times the average cosmic value and temperature of $10^5-10^7$K, forming the ``Warm-Hot Intergalactic Medium'' (WHIM, see e.g. \citealt[][]{1999ApJ...514....1C,2001ApJ...552..473D}, \citealt[][]{1996Natur.380..603B}, \citealt[and references therein]{do08, 2001ApJ...552..473D,2005MNRAS.360.1110V,2006MNRAS.370..656D}). Direct observations of the cosmic web are challenging, due to its extremely low mass density. Indications of possible detections of filamentary structures emerged from the analysis of soft X--ray \citep[e.g.][]{2003A&A...410..777F,2008A&A...482L..29W,2010ApJ...715..854N,2013ApJ...769...90N}, or Sunyaev-Zeldovich effect \citep[][]{2013A&A...550A.134P} data. However, very recently the first X-ray imaging of the terminal part of four filaments connected to the virial radius of cluster A2744 has been reported by \citet{2015Natur.528..105E} using XMM-Newton. The interest in the observational detection of filaments is growing rapidly since their evolution is less violent and complex than that of galaxy clusters, with adiabatic physics (besides gravity) dominating the gas dynamics. Therefore, filaments are expected to preserve many traces of the original environment in which the process of gravitational clustering started. They are also expected to retain memory of the initial magnetic seed fields of the Universe because they should not host strong dynamo amplification \citep[][]{ry08,donn09,va14mhd}. Cosmological simulations can now be used to produce observable predictions for different models of extragalactic magnetic fields, to be tested by radio observations. % Radio surveys at low frequencies (e.g. LOFAR, MWA and SKA) are expected to detect the brightest parts of the filamentary cosmic web if the magnetic field in the WHIM is amplified to the level of a few $\%$ of the thermal gas energy there \citep[][]{va15ska,va15radio}. Predicting the evolution of the magnetic fields in filaments is important because radio surveys might be able to detect the cosmic web in the redshift range $z \sim 0.1-0.2$, while especially at high radio frequencies it is impossible to sample the large angular scales ($\geq $ degrees) associated with the emission from the cosmic web in the local Universe \citep[][]{va15radio}.\\ In addition to galaxy-galaxy mergers \citep[e.g.][]{2003AJ....126.1183C,2006ApJ...645..986R,2007A&A...468...61D,2009ApJ...699L.178N} and the co-evolution of supermassive black holes \citep[e.g.][]{2005MNRAS.361..776S,2007MNRAS.382.1415S,2008ApJS..175..356H}, the relation between galaxies and their surrounding large-scale environment is key to understand the accretion of cold gas into galactic centres \citep[e.g.][]{2015MNRAS.454..637G}, the spin orientation of galaxies \citep[e.g.][]{2010MNRAS.405..274H} and their ellipticity \citep[e.g.][]{2015MNRAS.454.3341C}. On the other hand, other works suggest that the star formation rate is more dependent on the stellar mass and not on the environment \citep[e.g.][]{2012MNRAS.423.3679W}. More recently, \citet{alp15} observed that the large-scale environment of galaxies determines their stellar mass function, but has otherwise a modest impact on other galactic properties. An important motivation to study the filamentary structure of the cosmic web as well as the simulated galaxies in filaments is to understand which internal galactic properties are connected to the large-scale properties of filaments, and which are independent. We have studied in detail the properties of simulated dark-matter halos hosting galaxies in filaments, and compared them with those of real galaxies in the GAMA survey \citep[][]{driver09,driver11,liske15}.\\ Many high-resolution N-body cosmological simulations have been used to explore the properties of the cosmic web \citep[e.g.][]{2005MNRAS.359..272C,2007MNRAS.381...41H,2012MNRAS.421L.137L,2014MNRAS.441.2923C}. Most of these works, however, focus on the properties of the dark component, described via N-body algorithms, representing matter as a set of massive particles. In this work, we focus instead on the gas component, described by our code of choice, ENZO \citep{enzo13} as an Eulerian fluid on a computational mesh, which, for our purposes, has been kept at fixed resolution. We have exploited the fluid and thermodynamical properties of the gas to identify, extract and analyse the filamentary structures from the results of state-of-the art cosmological magneto-hydrodynamical simulations \citep[][]{va14mhd}. Our methodology, presented in all details in \citet{gh15} (hereafter Paper I), relies on the usage of the VisIt data visualization and analysis software \citep{Childs11visit:an}, exploiting a combination of its {\it Isovolume} and {\it Connected Components} algorithms. The adopted numerical and data processing methodologies are described in Sec. \ref{sec:nummethods}. In Sec. \ref{sec:statistics} and \ref{sec:galaxy}, we present the results obtained from several simulations of the $\Lambda$CDM cosmological model with density parameters $\Omega_0 = 1.0$, $\Omega_{\rm BM} = 0.0455$, $\Omega_{\rm DM} = 0.2265$ (BM and DM indicating the baryonic and the dark matter respectively), $\Omega_{\Lambda} = 0.728$, and a Hubble constant $H_0 = 70.2$km/sec/Mpc. In particular, Sec. \ref{sec:statistics} focuses on the study of the time evolution of the geometric, thermal and magnetic properties of filaments up to $z=1$. Sec. \ref{sec:galaxy} instead analyses the properties of the galactic halos hosted in filaments, in comparison to the results of the GAMA survey. The results are discussed in Sec. \ref{sec:discussion}, where we comment on the most important observational implications of our modelling of galaxies in filaments. The major achievements are finally summarized in Sec. \ref{sec:conclusions}.
\label{sec:conclusions} Cosmological filaments permeate the universe and account for a meaningful fraction of its mass. However, their detection is extremely challenging, due to their physical properties, which make them almost invisible at any wavelength. Therefore, it is crucial to find reliable approaches to identify a filament and characterize its main properties through easily detectable tracers. This would not only facilitate the filaments discovery process, but would also give essential hints to tune and optimize instruments and observations finalized to their detection. Our conclusions can be summarised as follows: \begin{itemize} \item In the redshift range explored ($0 \leq z \leq 1$) the thermal and magnetic properties evolve only mildly, e.g. the average temperature of the most massive objects decreases by a factor $\sim 2$ from z=0 to z=1 and also the mean and maximum magnetic fields evolve similarly. The length of filaments show a slow increase in this redshift range. This suggests that the properties of cosmic filaments are set already by $z \approx 1$, i.e. at a look-back time of $\sim 8$ Gyr. This is consistent with the fact that they are objects characterised by a very early formation epoch, and that their dynamical evolution is very slow. \item We find a small impact of the spatial/mass resolution on the global properties of the simulated population of filaments, with a factor $\sim 2$ in most cases. As resolution is increased, the thermalisation of filaments tends to decrease (due to the weakening of accretion shocks), the average length at a given mass is increased for the formation of more low density branches around filaments, and the magnetic field at a given mass is increased by the enhanced compression within filaments. \item The volume and density profiles of filaments show that the most important quantities (temperature, velocity, X-ray luminosity, magnetic fields) show a remarkable self-similar behaviour across the explored mass range. % \item Galaxies provide an effective proxy to characterize the geometric and physical properties of the hosting filaments, facilitating their detection through future observational campaigns. Future simulations including a more detailed description of galactic physics and higher spatial resolution will give even more accurate and comprehensive indications. \end{itemize}
16
7
1607.01406
Despite containing about a half of the total matter in the Universe, at most wavelengths the filamentary structure of the cosmic web is difficult to observe. In this work, we use large unigrid cosmological simulations to investigate how the geometrical, thermodynamical and magnetic properties of cosmological filaments vary with mass and redshift (z ≤ 1). We find that the average temperature, length, volume and magnetic field of filaments scales well with their total mass. This reflects the role of self-gravity in shaping their properties and enables statistical predictions of their observational properties based on their mass. We also focus on the properties of the simulated population of galaxy-sized haloes within filaments, and compare their properties to the results obtained from the spectroscopic GAMA survey. Simulated and observed filaments with the same length are found to contain an equal number of galaxies, with very similar distribution of masses. The total number of galaxies within each filament and the total/average stellar mass in galaxies can now be used to predict also the large-scale properties of the gas in the host filaments across tens or hundreds of Mpc in scale. These results are the first steps towards the future use of galaxy catalogues in order to select the best targets for observations of the warm-hot intergalactic medium.
false
[ "cosmological filaments", "filaments", "z ≤", "mass", "galaxy catalogues", "scale", "galaxies", "large unigrid cosmological simulations", "magnetic field", "statistical predictions", "most wavelengths", "masses", "the host filaments", "Mpc", "tens", "their observational properties", "hundreds", "their total mass", "the large-scale properties", "redshift" ]
12.084696
5.071469
-1
12406440
[ "Popovas, A.", "Jørgensen, U. G." ]
2016A&A...595A.130P
[ "Partition functions. I. Improved partition functions and thermodynamic quantities for normal, equilibrium, and ortho and para molecular hydrogen" ]
18
[ "Niels Bohr Institute &amp; Centre for Star and Planet Formation, University of Copenhagen, Øster Voldgade 5-7, 1350, Copenhagen K, Denmark", "Niels Bohr Institute &amp; Centre for Star and Planet Formation, University of Copenhagen, Øster Voldgade 5-7, 1350, Copenhagen K, Denmark" ]
[ "2018ASPC..515..145Z", "2019A&A...630A..82L", "2019A&A...632A.118S", "2019JPCA..123.9905G", "2019MNRAS.488.1846N", "2019MNRAS.489.4890B", "2020AcAC.1104...28G", "2020ApPhB.126...83R", "2021JPCA..125..795B", "2021JPCA..125.9226Z", "2021MNRAS.502.3394N", "2022A&A...664A..91P", "2022InJPh..96.2243G", "2022PhRvA.105a2425S", "2023JPCA..127.6842V", "2023MNRAS.518.2320D", "2023arXiv231202322J", "2024ApJ...965L..13A" ]
[ "astronomy", "physics" ]
9
[ "miscellaneous", "molecular data", "astrochemistry", "equation of state", "Physics - Chemical Physics", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1932PhRv...41..721D", "1950msms.book.....H", "1965MNRAS.129..345V", "1966PDAO...13....1T", "1970JChPh..52.2575T", "1981ApJS...45..621I", "1984A&A...130..202B", "1984ApJS...56..193S", "1984CaJPh..62.1639D", "1985A&A...148...93R", "1985JChPh..83.3034J", "1987A&A...182..348I", "1987JChPh..87.6908S", "1988JChPh..88..356M", "1990JChPh..93.2801M", "1990JMoSp.142..205G", "1993JChPh..99.3153S", "1993JMoSp.157..512A", "1994ApJS...95..535G", "1994CaJPh..72..856A", "1994JChPh.101.7692Y", "1994PhRvA..50.4618R", "2000JMoSt.517..407G", "2003JQSRT..82..401F", "2006JMoSp.238..118W", "2007ApJ...656L..89B", "2007JChPh.126n4303G", "2010MolPh.108..827B", "2011Icar..215..391L", "2012arXiv1201.5885T", "2013JMoSp.293....1G", "2013JMoSp.293...11G", "2013JMoSp.293...19G" ]
[ "10.1051/0004-6361/201527209", "10.48550/arXiv.1607.04479" ]
1607
1607.04479_arXiv.txt
The total internal partition function, $Q_{tot}(T)$, is used to determine how atoms and molecules in thermodynamic equilibrium are distributed among the various energy states at particular temperatures. It is the statistical sum over all the Boltzmann factors for all the bound levels. If the particle is not isolated, there is an occupation probability between 0 and 1 for each level depending on interactions with its neighbours. Together with other thermodynamic quantities, partition functions are used in many astrophysical problems, including equation of state, radiative transfer, dissociation equilibrium, evaluating line intensities from spectra, and correction of line intensities to temperatures other than given in standard atlases. Owing to the importance of $Q_{tot}(T)$, a number of studies were conducted throughout the past several decades to obtain more accurate values and present them in a convenient way. It is essential that a standard coherent set of $Q_{tot}(T)$ is being used for any meaningful astrophysical conclusions from calculations of different atmospheric models and their comparisons. Unfortunately, today we face a completely different situation. We have noted that most studies give more or less different results, sometimes the differences are small, but sometimes they are quite dramatic, for instance when different studies used different conventions to treat nuclear spin states (and later do not strictly specify these) or different approximations, cut-offs, etc. Furthermore, different methods of calculating the $Q_{tot}(T)$ are implemented in different codes, and it is not always clear which methods in particular are used. Naturally, differences in $Q_{tot}(T)$ values and hence in their derivatives (internal energy, specific heat, entropy, free energy) result in differences in the physical structure of computed model atmosphere even when line lists and input physical quantities (e.g. $T_{eff}$, log$g$, and metallicities) are identical. In the subsequent sections we review how $Q_{tot}(T)$ is calculated, comment on which simplifications, approximations, and constraints are used in a number of studies, and show how they compare to each other. We also argue against using molecular constants to calculate the partition functions. The molecular constants are rooted in the semi-classical idea that the vibrational-rotational eigenvalues can be expressed in terms of a modified classical oscillator-rotor analogue. This erroneous concept has severe challenges at the highest energy levels that are not avoided by instead making a simple summation of experimentally determined energy levels, simply due to the necessity of naming the experimental levels by use of assigned quantum numbers. Instead, we report Dunham\cite{Dunham1932} coefficients for the ground electronic state and a number of excited electronic states, as well as resulting partition functions and thermodynamic quantities. The Dumham coefficients are not rooted in the classical picture, but still use quantum number assignments of the energy levels, and this approach is also bound to the same challenges of defining the upper energy levels in the summation, as is the summation using molecular constants (and/or pure experimental data). No published studies have solved this challenge yet, but we quantify the uncertainty it implies on the resulting values of the chemical equilibrium partition function.
We have investigated the shortcomings of various simplifications that are used to calculate $Q_{int}$ and applied our analyses to calculate the time-independent partition function of normal, ortho and para, and equilibrium molecular hydrogen. We also calculated $E_{int}/RT$, $H-H(0)$, $S$, $C_p$, $C_v$, $-[G-H(0)]/T,$ and $\gamma$. Our calculated values of thermodynamic properties for ortho and para, equilibrium, and normal $H_2$ are presented in Tables \ref{tab:result-thermo-Qnorm}, \ref{tab:result-thermo-Qeq}, \ref{tab:result-thermo-Qortho}, and \ref{tab:result-thermo-Qpara}, rounded to three significant digits. Full datasets in 1 K temperature steps\footnote{Data with smaller temperature steps or beyond the given temperature range can be requested personally from the corresponding author} can be retrieved online from \cite{mydata}. The partition functions and thermodynamic data presented in this work are more accurate than previously available data from the literature and cover a more complete temperature range than any previous study in the literature. Determined Dunham coefficients for a number of electronic states of $\rm H_2$ are also reported.\\ In future work we plan to use the method we described here to calculate partition functions and thermodynamic quantities for other astrophysically important molecular species as well. \onltab{ \begin{table*}[h] \caption{Thermodynamic properties of equilibrium $\rm H_2$.} \begin{tabular}{ccccccccc} % \hline $T$ & $Q_{int}$ & $E_{int}/RT$ & $H-H(0)$ & $S$ & $C_p$ & $C_v$ & $-[G-H(0)]/T$ &$\gamma$\\ [K] & & [J] & [J mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J $K^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] \\ \hline 5.0 & 0.25 & 0.0 & 103.931 & 21.095 & 20.786 & 0.0 & -40.026 & 1.667\\ 10.0 & 0.25 & 0.0 & 207.862 & 35.503 & 20.787 & 0.001 & 64.026 & 1.667\\ 15.0 & 0.25 & 0.001 & 311.94 & 43.942 & 20.898 & 0.112 & 222.473 & 1.666\\ 20.0 & 0.25 & 0.015 & 418.257 & 50.053 & 21.864 & 1.078 & 416.507 & 1.66\\ 25.0 & 0.252 & 0.066 & 533.473 & 55.184 & 24.518 & 3.732 & 638.256 & 1.638\\ 30.0 & 0.258 & 0.169 & 665.765 & 59.996 & 28.537 & 7.751 & 884.693 & 1.599\\ 35.0 & 0.267 & 0.315 & 819.085 & 64.715 & 32.706 & 11.92 & 1154.933 & 1.551\\ 40.0 & 0.282 & 0.48 & 991.15 & 69.305 & 35.893 & 15.107 & 1448.489 & 1.505\\ 45.0 & 0.301 & 0.642 & 1175.533 & 73.647 & 37.61 & 16.824 & 1764.426 & 1.467\\ 50.0 & 0.324 & 0.783 & 1364.975 & 77.638 & 37.97 & 17.184 & 2101.227 & 1.438\\ 60.0 & 0.382 & 0.984 & 1737.824 & 84.441 & 36.242 & 15.456 & 2829.782 & 1.403\\ 70.0 & 0.448 & 1.085 & 2086.766 & 89.826 & 33.537 & 12.75 & 3619.013 & 1.387\\ 80.0 & 0.519 & 1.123 & 2409.737 & 94.142 & 31.153 & 10.367 & 4456.451 & 1.381\\ 90.0 & 0.593 & 1.124 & 2711.881 & 97.703 & 29.377 & 8.591 & 5333.052 & 1.381\\ 100.0 & 0.667 & 1.107 & 2999.115 & 100.73 & 28.151 & 7.365 & 6242.444 & 1.384\\ 110.0 & 0.74 & 1.082 & 3276.322 & 103.373 & 27.351 & 6.565 & 7180.087 & 1.387\\ 120.0 & 0.813 & 1.055 & 3547.19 & 105.73 & 26.867 & 6.081 & 8142.661 & 1.391\\ 130.0 & 0.883 & 1.029 & 3814.429 & 107.869 & 26.613 & 5.827 & 9127.671 & 1.395\\ 140.0 & 0.952 & 1.005 & 4080.006 & 109.837 & 26.526 & 5.739 & 10133.185 & 1.399\\ 150.0 & 1.02 & 0.984 & 4345.331 & 111.668 & 26.556 & 5.77 & 11157.669 & 1.403\\ 200.0 & 1.341 & 0.923 & 5692.821 & 119.414 & 27.448 & 6.662 & 16527.172 & 1.413\\ 150.0 & 1.02 & 0.984 & 4345.331 & 111.668 & 26.556 & 5.77 & 11157.669 & 1.403\\ 298.15 & 1.931 & 0.916 & 8467.176 & 130.682 & 28.836 & 8.05 & 28016.705 & 1.414\\ 300.0 & 1.942 & 0.916 & 8520.534 & 130.86 & 28.849 & 8.063 & 28243.25 & 1.414\\ 350.0 & 2.238 & 0.926 & 9969.563 & 135.327 & 29.081 & 8.294 & 34484.941 & 1.412\\ 400.0 & 2.534 & 0.936 & 11426.436 & 139.218 & 29.181 & 8.395 & 40934.958 & 1.411\\ 450.0 & 2.831 & 0.944 & 12886.795 & 142.658 & 29.229 & 8.442 & 47567.798 & 1.409\\ 500.0 & 3.128 & 0.952 & 14349.02 & 145.739 & 29.259 & 8.473 & 54363.344 & 1.408\\ 600.0 & 3.725 & 0.963 & 17278.057 & 151.079 & 29.326 & 8.54 & 68380.83 & 1.406\\ 700.0 & 4.324 & 0.973 & 20215.836 & 155.608 & 29.44 & 8.654 & 82889.4 & 1.404\\ 800.0 & 4.928 & 0.983 & 23168.335 & 159.55 & 29.623 & 8.837 & 97820.014 & 1.403\\ 900.0 & 5.536 & 0.994 & 26142.864 & 163.053 & 29.88 & 9.094 & 113121.903 & 1.401\\ 1000.0 & 6.151 & 1.006 & 29146.545 & 166.217 & 30.204 & 9.418 & 128756.479 & 1.399\\ 1100.0 & 6.774 & 1.019 & 32185.337 & 169.113 & 30.579 & 9.793 & 144693.583 & 1.397\\ 1200.0 & 7.406 & 1.034 & 35263.602 & 171.792 & 30.991 & 10.204 & 160909.037 & 1.395\\ 1300.0 & 8.051 & 1.051 & 38384.094 & 174.289 & 31.421 & 10.635 & 177383.01 & 1.392\\ 1400.0 & 8.709 & 1.069 & 41548.164 & 176.634 & 31.86 & 11.074 & 194098.888 & 1.389\\ 1500.0 & 9.382 & 1.089 & 44756.054 & 178.847 & 32.297 & 11.511 & 211042.494 & 1.386\\ 2000.0 & 13.015 & 1.193 & 61416.861 & 188.419 & 34.278 & 13.492 & 298792.538 & 1.371\\ 2500.0 & 17.181 & 1.299 & 78962.143 & 196.243 & 35.838 & 15.052 & 390859.016 & 1.357\\ 3000.0 & 21.964 & 1.397 & 97201.242 & 202.89 & 37.077 & 16.291 & 486526.148 & 1.345\\ 3500.0 & 27.428 & 1.486 & 116007.0 & 208.686 & 38.121 & 17.335 & 585293.496 & 1.335\\ 4000.0 & 33.632 & 1.568 & 135302.809 & 213.838 & 39.043 & 18.257 & 686790.762 & 1.326\\ 4500.0 & 40.634 & 1.644 & 155031.8 & 218.484 & 39.852 & 19.066 & 790732.925 & 1.318\\ 5000.0 & 48.492 & 1.713 & 175128.639 & 222.719 & 40.508 & 19.721 & 896892.151 & 1.311\\ 6000.0 & 66.988 & 1.831 & 216042.905 & 230.176 & 41.156 & 20.37 & 1.115128382e6 & 1.3\\ 7000.0 & 89.448 & 1.917 & 257096.393 & 236.506 & 40.779 & 19.993 & 1.340240895e6 & 1.293\\ 8000.0 & 115.989 & 1.97 & 297322.673 & 241.879 & 39.553 & 18.767 & 1.571192049e6 & 1.288\\ 9000.0 & 146.498 & 1.991 & 336046.785 & 246.442 & 37.836 & 17.049 & 1.807099542e6 & 1.286\\ 10000.0 & 180.675 & 1.986 & 372961.362 & 250.333 & 35.979 & 15.193 & 2.047222929e6 & 1.287\\ 11000.0 & 218.099 & 1.962 & 408052.021 & 253.679 & 34.211 & 13.425 & 2.29095485e6 & 1.289\\ 12000.0 & 258.301 & 1.925 & 441490.549 & 256.589 & 32.676 & 11.89 & 2.537806546e6 & 1.292\\ 13000.0 & 300.818 & 1.881 & 473549.464 & 259.156 & 31.484 & 10.698 & 2.787389971e6 & 1.296\\ 14000.0 & 345.231 & 1.834 & 504545.723 & 261.453 & 30.59 & 9.804 & 3.039399948e6 & 1.3\\ 15000.0 & 391.19 & 1.788 & 534806.189 & 263.541 & 29.996 & 9.21 & 3.293598228e6 & 1.304\\ 16000.0 & 438.423 & 1.744 & 564646.464 & 265.467 & 29.722 & 8.936 & 3.549799968e6 & 1.308\\ 17000.0 & 486.741 & 1.705 & 594357.51 & 267.269 & 29.744 & 8.958 & 3.80786255e6 & 1.312\\ 18000.0 & 536.03 & 1.671 & 624196.974 & 268.974 & 29.964 & 9.178 & 4.067676383e6 & 1.315\\ 19000.0 & 586.244 & 1.642 & 654383.656 & 270.606 & 30.448 & 9.662 & 4.329157371e6 & 1.318\\ 20000.0 & 637.389 & 1.62 & 685094.498 & 272.181 & 39.3 & 18.514 & 4.59224074e6 & 1.321\\ \end{tabular} \label{tab:result-thermo-Qeq} \end{table*}} \onltab{ \begin{table*}[h] \caption{Thermodynamic properties of normal $\rm H_2$.} \begin{tabular}{ccccccccc} % \hline $T$ & $Q_{int}$ & $E_{int}/RT$ & $H-H(0)$ & $S$ & $C_p$ & $C_v$ & $-[G-H(0)]/T$ &$\gamma$\\ [K] & & [J] & [J mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J $K^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] \\ \hline 5.0 & 2.28 & 0.0 & 103.931 & 39.472 & 20.786 & 0.0 & 51.859 & 1.667\\ 10.0 & 2.28 & 0.0 & 207.861 & 53.88 & 20.786 & 0.0 & 247.797 & 1.667\\ 15.0 & 2.28 & 0.0 & 311.792 & 62.308 & 20.786 & -0.0 & 498.116 & 1.667\\ 20.0 & 2.28 & 0.0 & 415.723 & 68.288 & 20.786 & 0.0 & 783.751 & 1.667\\ 25.0 & 2.28 & 0.0 & 519.654 & 72.926 & 20.786 & 0.0 & 1095.646 & 1.667\\ 30.0 & 2.28 & 0.0 & 623.585 & 76.716 & 20.786 & 0.0 & 1428.468 & 1.667\\ 35.0 & 2.28 & 0.0 & 727.518 & 79.92 & 20.787 & 0.001 & 1778.693 & 1.667\\ 40.0 & 2.28 & 0.0 & 831.461 & 82.696 & 20.791 & 0.005 & 2143.818 & 1.667\\ 45.0 & 2.28 & 0.0 & 935.441 & 85.146 & 20.802 & 0.016 & 2521.971 & 1.667\\ 50.0 & 2.28 & 0.0 & 1039.506 & 87.339 & 20.827 & 0.041 & 2911.705 & 1.666\\ 60.0 & 2.28 & 0.002 & 1248.259 & 91.144 & 20.941 & 0.155 & 3721.536 & 1.666\\ 70.0 & 2.281 & 0.006 & 1458.731 & 94.388 & 21.175 & 0.389 & 4566.442 & 1.664\\ 80.0 & 2.284 & 0.014 & 1672.186 & 97.238 & 21.537 & 0.751 & 5441.706 & 1.661\\ 90.0 & 2.29 & 0.026 & 1889.847 & 99.801 & 22.011 & 1.225 & 6343.964 & 1.656\\ 100.0 & 2.298 & 0.041 & 2112.669 & 102.148 & 22.563 & 1.777 & 7270.725 & 1.649\\ 110.0 & 2.309 & 0.06 & 2341.243 & 104.326 & 23.155 & 2.369 & 8220.08 & 1.641\\ 120.0 & 2.323 & 0.082 & 2575.795 & 106.367 & 23.754 & 2.968 & 9190.508 & 1.632\\ 130.0 & 2.34 & 0.106 & 2816.26 & 108.291 & 24.334 & 3.548 & 10180.745 & 1.623\\ 140.0 & 2.361 & 0.131 & 3062.367 & 110.115 & 24.881 & 4.094 & 11189.711 & 1.613\\ 150.0 & 2.384 & 0.157 & 3313.73 & 111.849 & 25.385 & 4.599 & 12216.458 & 1.603\\ 200.0 & 2.54 & 0.287 & 4634.245 & 119.434 & 27.269 & 6.483 & 17589.625 & 1.56\\ 150.0 & 2.384 & 0.157 & 3313.73 & 111.849 & 25.385 & 4.599 & 12216.458 & 1.603\\ 298.15 & 2.964 & 0.487 & 7404.367 & 130.682 & 28.834 & 8.047 & 29079.567 & 1.503\\ 300.0 & 2.973 & 0.49 & 7457.721 & 130.861 & 28.846 & 8.06 & 29306.112 & 1.503\\ 350.0 & 3.224 & 0.561 & 8906.702 & 135.327 & 29.08 & 8.294 & 35547.806 & 1.485\\ 400.0 & 3.488 & 0.616 & 10363.571 & 139.218 & 29.181 & 8.395 & 41997.824 & 1.473\\ 450.0 & 3.761 & 0.66 & 11823.929 & 142.658 & 29.229 & 8.442 & 48630.664 & 1.463\\ 500.0 & 4.039 & 0.696 & 13286.154 & 145.739 & 29.259 & 8.473 & 55426.21 & 1.455\\ 600.0 & 4.609 & 0.75 & 16215.191 & 151.079 & 29.326 & 8.54 & 69443.696 & 1.444\\ 700.0 & 5.191 & 0.791 & 19152.97 & 155.608 & 29.44 & 8.654 & 83952.266 & 1.437\\ 800.0 & 5.782 & 0.823 & 22105.469 & 159.55 & 29.623 & 8.837 & 98882.88 & 1.43\\ 900.0 & 6.381 & 0.852 & 25079.998 & 163.053 & 29.88 & 9.094 & 114184.769 & 1.425\\ 1000.0 & 6.989 & 0.878 & 28083.68 & 166.217 & 30.204 & 9.418 & 129819.345 & 1.421\\ 1100.0 & 7.608 & 0.903 & 31122.471 & 169.113 & 30.579 & 9.793 & 145756.449 & 1.416\\ 1200.0 & 8.239 & 0.928 & 34200.737 & 171.792 & 30.99 & 10.204 & 161971.903 & 1.412\\ 1300.0 & 8.883 & 0.953 & 37321.228 & 174.289 & 31.422 & 10.635 & 178445.876 & 1.408\\ 1400.0 & 9.542 & 0.978 & 40485.298 & 176.634 & 31.86 & 11.074 & 195161.754 & 1.404\\ 1500.0 & 10.217 & 1.003 & 43693.188 & 178.847 & 32.297 & 11.51 & 212105.359 & 1.399\\ 2000.0 & 13.874 & 1.129 & 60353.995 & 188.419 & 34.278 & 13.492 & 299855.404 & 1.38\\ 2500.0 & 18.083 & 1.248 & 77899.277 & 196.243 & 35.838 & 15.052 & 391921.882 & 1.364\\ 3000.0 & 22.92 & 1.354 & 96138.376 & 202.89 & 37.077 & 16.291 & 487589.013 & 1.35\\ 3500.0 & 28.448 & 1.45 & 114944.134 & 208.686 & 38.122 & 17.336 & 586356.362 & 1.339\\ 4000.0 & 34.724 & 1.536 & 134239.943 & 213.838 & 39.043 & 18.257 & 687853.627 & 1.329\\ 4500.0 & 41.805 & 1.615 & 153968.934 & 218.484 & 39.852 & 19.065 & 791795.791 & 1.321\\ 5000.0 & 49.748 & 1.687 & 174065.773 & 222.719 & 40.508 & 19.722 & 897955.017 & 1.314\\ 6000.0 & 68.431 & 1.809 & 214980.038 & 230.176 & 41.156 & 20.37 & 1.116191248e6 & 1.302\\ 7000.0 & 91.097 & 1.899 & 256033.52 & 236.505 & 40.784 & 19.998 & 1.34130376e6 & 1.294\\ 8000.0 & 117.857 & 1.954 & 296259.779 & 241.879 & 39.56 & 18.774 & 1.572254912e6 & 1.29\\ 9000.0 & 148.594 & 1.977 & 334983.848 & 246.442 & 37.84 & 17.054 & 1.808162398e6 & 1.288\\ 10000.0 & 182.999 & 1.973 & 371898.356 & 250.333 & 35.983 & 15.197 & 2.048285773e6 & 1.288\\ 11000.0 & 220.648 & 1.95 & 406988.917 & 253.679 & 34.219 & 13.432 & 2.292017674e6 & 1.29\\ 12000.0 & 261.068 & 1.914 & 440427.327 & 256.589 & 32.717 & 11.931 & 2.538869339e6 & 1.293\\ 13000.0 & 303.791 & 1.871 & 472486.104 & 259.156 & 31.482 & 10.696 & 2.788452722e6 & 1.297\\ 14000.0 & 348.397 & 1.825 & 503482.216 & 261.453 & 30.569 & 9.782 & 3.040462648e6 & 1.301\\ 15000.0 & 394.537 & 1.78 & 533742.527 & 263.541 & 30.004 & 9.217 & 3.294660864e6 & 1.305\\ 16000.0 & 441.939 & 1.736 & 563582.645 & 265.467 & 29.741 & 8.955 & 3.550862531e6 & 1.309\\ 17000.0 & 490.413 & 1.697 & 593293.539 & 267.268 & 29.736 & 8.95 & 3.808925029e6 & 1.313\\ 18000.0 & 539.849 & 1.664 & 623132.856 & 268.974 & 29.995 & 9.209 & 4.06873877e6 & 1.316\\ 19000.0 & 590.199 & 1.636 & 653319.399 & 270.606 & 30.443 & 9.657 & 4.330219658e6 & 1.319\\ 20000.0 & 641.473 & 1.613 & 684030.114 & 272.181 & 39.297 & 18.511 & 4.593302921e6 & 1.321\\ \end{tabular} \label{tab:result-thermo-Qnorm} \end{table*}} \onltab{ \begin{table*}[h] \caption{Thermodynamic properties of ortho $\rm H_2$.} \begin{tabular}{ccccccccc} % \hline $T$ & $Q_{int}$ & $E_{int}/RT$ & $H-H(0)$ & $S$ & $C_p$ & $C_v$ & $-[G-H(0)]/T$ &$\gamma$\\ [K] & & [J] & [J mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J $K^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] \\ \hline 5.0 & 3.0 & 0.0 & 103.931 & 41.756 & 20.786 & 0.0 & 63.277 & 1.667\\ 10.0 & 3.0 & 0.0 & 207.861 & 56.164 & 20.786 & 0.0 & 270.633 & 1.667\\ 15.0 & 3.0 & 0.0 & 311.792 & 64.592 & 20.786 & 0.0 & 532.37 & 1.667\\ 20.0 & 3.0 & 0.0 & 415.723 & 70.572 & 20.786 & 0.0 & 829.423 & 1.667\\ 25.0 & 3.0 & 0.0 & 519.654 & 75.21 & 20.786 & 0.0 & 1152.736 & 1.667\\ 30.0 & 3.0 & 0.0 & 623.584 & 79.0 & 20.786 & 0.0 & 1496.976 & 1.667\\ 35.0 & 3.0 & 0.0 & 727.515 & 82.204 & 20.786 & 0.0 & 1858.619 & 1.667\\ 40.0 & 3.0 & 0.0 & 831.446 & 84.98 & 20.786 & 0.0 & 2235.16 & 1.667\\ 45.0 & 3.0 & 0.0 & 935.377 & 87.428 & 20.786 & 0.0 & 2624.727 & 1.667\\ 50.0 & 3.0 & 0.0 & 1039.308 & 89.618 & 20.786 & 0.0 & 3025.865 & 1.667\\ 60.0 & 3.0 & 0.0 & 1247.182 & 93.408 & 20.789 & 0.003 & 3858.425 & 1.667\\ 70.0 & 3.0 & 0.0 & 1455.125 & 96.613 & 20.802 & 0.016 & 4725.796 & 1.667\\ 80.0 & 3.0 & 0.001 & 1663.318 & 99.393 & 20.842 & 0.056 & 5622.99 & 1.666\\ 90.0 & 3.001 & 0.002 & 1872.131 & 101.853 & 20.93 & 0.144 & 6546.308 & 1.666\\ 100.0 & 3.002 & 0.004 & 2082.135 & 104.065 & 21.083 & 0.297 & 7492.933 & 1.665\\ 110.0 & 3.003 & 0.008 & 2294.057 & 106.085 & 21.315 & 0.529 & 8460.68 & 1.663\\ 120.0 & 3.006 & 0.014 & 2508.698 & 107.952 & 21.627 & 0.84 & 9447.833 & 1.66\\ 130.0 & 3.011 & 0.023 & 2726.839 & 109.698 & 22.013 & 1.227 & 10453.031 & 1.657\\ 140.0 & 3.017 & 0.034 & 2949.171 & 111.346 & 22.462 & 1.676 & 11475.178 & 1.652\\ 150.0 & 3.025 & 0.047 & 3176.238 & 112.912 & 22.958 & 2.171 & 12513.382 & 1.647\\ 200.0 & 3.103 & 0.14 & 4390.077 & 119.878 & 25.561 & 4.775 & 17922.591 & 1.61\\ 150.0 & 3.025 & 0.047 & 3176.238 & 112.912 & 22.958 & 2.171 & 12513.382 & 1.647\\ 298.15 & 3.416 & 0.352 & 7069.282 & 130.738 & 28.461 & 7.675 & 29431.377 & 1.54\\ 300.0 & 3.424 & 0.355 & 7121.959 & 130.914 & 28.486 & 7.7 & 29658.024 & 1.539\\ 350.0 & 3.641 & 0.441 & 8559.265 & 135.345 & 28.941 & 8.155 & 35901.334 & 1.515\\ 400.0 & 3.88 & 0.51 & 10011.739 & 139.223 & 29.13 & 8.344 & 42351.87 & 1.497\\ 450.0 & 4.134 & 0.566 & 11470.509 & 142.66 & 29.21 & 8.424 & 48984.875 & 1.484\\ 500.0 & 4.399 & 0.611 & 12932.172 & 145.74 & 29.253 & 8.467 & 55780.474 & 1.474\\ 600.0 & 4.948 & 0.679 & 15860.941 & 151.079 & 29.325 & 8.539 & 69797.982 & 1.459\\ 700.0 & 5.517 & 0.73 & 18798.687 & 155.608 & 29.44 & 8.653 & 84306.554 & 1.448\\ 800.0 & 6.098 & 0.77 & 21751.181 & 159.55 & 29.623 & 8.837 & 99237.169 & 1.441\\ 900.0 & 6.69 & 0.804 & 24725.709 & 163.053 & 29.88 & 9.094 & 114539.057 & 1.434\\ 1000.0 & 7.294 & 0.835 & 27729.391 & 166.217 & 30.204 & 9.418 & 130173.634 & 1.428\\ 1100.0 & 7.909 & 0.864 & 30768.183 & 169.113 & 30.579 & 9.793 & 146110.737 & 1.423\\ 1200.0 & 8.537 & 0.892 & 33846.448 & 171.792 & 30.99 & 10.204 & 162326.192 & 1.418\\ 1300.0 & 9.179 & 0.92 & 36966.94 & 174.289 & 31.422 & 10.635 & 178800.165 & 1.413\\ 1400.0 & 9.836 & 0.948 & 40131.01 & 176.634 & 31.86 & 11.074 & 195516.042 & 1.409\\ 1500.0 & 10.511 & 0.975 & 43338.899 & 178.847 & 32.297 & 11.51 & 212459.648 & 1.404\\ 2000.0 & 14.173 & 1.108 & 59999.706 & 188.419 & 34.278 & 13.492 & 300209.693 & 1.383\\ 2500.0 & 18.393 & 1.231 & 77544.988 & 196.243 & 35.838 & 15.052 & 392276.17 & 1.366\\ 3000.0 & 23.248 & 1.34 & 95784.081 & 202.89 & 37.077 & 16.291 & 487943.302 & 1.352\\ 3500.0 & 28.797 & 1.438 & 114589.782 & 208.686 & 38.122 & 17.335 & 586710.645 & 1.34\\ 4000.0 & 35.096 & 1.526 & 133885.328 & 213.838 & 39.042 & 18.256 & 688207.888 & 1.331\\ 4500.0 & 42.202 & 1.606 & 153613.51 & 218.484 & 39.849 & 19.063 & 792149.966 & 1.322\\ 5000.0 & 50.173 & 1.678 & 173708.508 & 222.718 & 40.503 & 19.717 & 898308.966 & 1.315\\ 6000.0 & 68.916 & 1.802 & 214614.043 & 230.174 & 41.144 & 20.357 & 1.116543844e6 & 1.303\\ 7000.0 & 91.645 & 1.893 & 255650.45 & 236.501 & 40.76 & 19.974 & 1.34165289e6 & 1.295\\ 8000.0 & 118.466 & 1.948 & 295852.244 & 241.871 & 39.533 & 18.746 & 1.572597596e6 & 1.29\\ 9000.0 & 149.256 & 1.971 & 334547.337 & 246.43 & 37.81 & 17.024 & 1.808495266e6 & 1.288\\ 10000.0 & 183.704 & 1.967 & 371431.338 & 250.318 & 35.952 & 15.166 & 2.048605493e6 & 1.288\\ 11000.0 & 221.382 & 1.945 & 406492.024 & 253.661 & 34.191 & 13.405 & 2.292321206e6 & 1.29\\ 12000.0 & 261.814 & 1.909 & 439902.437 & 256.569 & 32.688 & 11.902 & 2.53915404e6 & 1.293\\ 13000.0 & 304.533 & 1.866 & 471935.664 & 259.134 & 31.454 & 10.668 & 2.788716353e6 & 1.297\\ 14000.0 & 349.119 & 1.82 & 502908.834 & 261.43 & 30.545 & 9.758 & 3.040703339e6 & 1.301\\ 15000.0 & 395.221 & 1.775 & 533148.787 & 263.516 & 29.974 & 9.188 & 3.294877066e6 & 1.305\\ 16000.0 & 442.572 & 1.732 & 562971.041 & 265.441 & 29.735 & 8.949 & 3.551052967e6 & 1.309\\ 17000.0 & 490.982 & 1.693 & 592666.488 & 267.241 & 29.721 & 8.935 & 3.809088657e6 & 1.313\\ 18000.0 & 540.34 & 1.659 & 622492.742 & 268.946 & 29.99 & 9.204 & 4.068874749e6 & 1.317\\ 19000.0 & 590.602 & 1.631 & 652668.607 & 270.577 & 30.434 & 9.648 & 4.330327328e6 & 1.319\\ 20000.0 & 641.778 & 1.61 & 683371.061 & 272.152 & 39.29 & 18.504 & 4.59338178e6 & 1.322\\ \end{tabular} \label{tab:result-thermo-Qortho} \end{table*}} \onltab{ \begin{table*}[h] \caption{Thermodynamic properties of para $\rm H_2$.} \begin{tabular}{ccccccccc} % \hline $T$ & $Q_{int}$ & $E_{int}/RT$ & $H-H(0)$ & $S$ & $C_p$ & $C_v$ & $-[G-H(0)]/T$ &$\gamma$\\ [K] & & [J] & [J mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J $K^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] & [J K$^{-1}$ mol$^{-1}$] \\ \hline 5.0 & 1.0 & 0.0 & 103.931 & 32.622 & 20.786 & 0.0 & 17.605 & 1.667\\ 10.0 & 1.0 & 0.0 & 207.861 & 47.03 & 20.786 & 0.0 & 179.289 & 1.667\\ 15.0 & 1.0 & 0.0 & 311.792 & 55.458 & 20.786 & 0.0 & 395.355 & 1.667\\ 20.0 & 1.0 & 0.0 & 415.723 & 61.437 & 20.786 & 0.0 & 646.735 & 1.667\\ 25.0 & 1.0 & 0.0 & 519.654 & 66.076 & 20.786 & 0.0 & 924.377 & 1.667\\ 30.0 & 1.0 & 0.0 & 623.585 & 69.865 & 20.787 & 0.001 & 1222.945 & 1.667\\ 35.0 & 1.0 & 0.0 & 727.525 & 73.07 & 20.79 & 0.004 & 1538.917 & 1.667\\ 40.0 & 1.0 & 0.0 & 831.508 & 75.847 & 20.806 & 0.02 & 1869.79 & 1.667\\ 45.0 & 1.0 & 0.001 & 935.632 & 78.3 & 20.85 & 0.064 & 2213.703 & 1.666\\ 50.0 & 1.0 & 0.002 & 1040.099 & 80.501 & 20.947 & 0.161 & 2569.224 & 1.666\\ 60.0 & 1.001 & 0.009 & 1251.492 & 84.354 & 21.398 & 0.612 & 3310.871 & 1.663\\ 70.0 & 1.003 & 0.025 & 1469.549 & 87.713 & 22.291 & 1.505 & 4088.379 & 1.656\\ 80.0 & 1.009 & 0.054 & 1698.789 & 90.773 & 23.621 & 2.835 & 4897.857 & 1.644\\ 90.0 & 1.017 & 0.097 & 1942.994 & 93.647 & 25.255 & 4.469 & 5736.931 & 1.626\\ 100.0 & 1.031 & 0.151 & 2204.272 & 96.398 & 27.003 & 6.217 & 6604.1 & 1.606\\ 110.0 & 1.049 & 0.215 & 2482.799 & 99.052 & 28.677 & 7.891 & 7498.282 & 1.583\\ 120.0 & 1.071 & 0.283 & 2777.086 & 101.611 & 30.136 & 9.35 & 8418.531 & 1.561\\ 130.0 & 1.099 & 0.354 & 3084.521 & 104.071 & 31.298 & 10.511 & 9363.887 & 1.539\\ 140.0 & 1.131 & 0.423 & 3401.955 & 106.423 & 32.136 & 11.349 & 10333.31 & 1.52\\ 150.0 & 1.167 & 0.488 & 3726.206 & 108.66 & 32.666 & 11.88 & 11325.683 & 1.503\\ 200.0 & 1.393 & 0.727 & 5366.75 & 118.102 & 32.393 & 11.607 & 16590.726 & 1.449\\ 150.0 & 1.167 & 0.488 & 3726.206 & 108.66 & 32.666 & 11.88 & 11325.683 & 1.503\\ 298.15 & 1.937 & 0.892 & 8409.623 & 130.514 & 29.951 & 9.165 & 28024.137 & 1.418\\ 300.0 & 1.947 & 0.894 & 8465.01 & 130.699 & 29.927 & 9.141 & 28250.377 & 1.418\\ 350.0 & 2.24 & 0.919 & 9949.013 & 135.275 & 29.497 & 8.711 & 34487.222 & 1.413\\ 400.0 & 2.535 & 0.933 & 11419.066 & 139.201 & 29.333 & 8.547 & 40935.684 & 1.411\\ 450.0 & 2.831 & 0.944 & 12884.19 & 142.653 & 29.283 & 8.497 & 47568.031 & 1.409\\ 500.0 & 3.128 & 0.951 & 14348.103 & 145.738 & 29.278 & 8.492 & 54363.419 & 1.408\\ 600.0 & 3.725 & 0.963 & 17277.941 & 151.079 & 29.329 & 8.542 & 68380.839 & 1.406\\ 700.0 & 4.324 & 0.973 & 20215.82 & 155.608 & 29.44 & 8.654 & 82889.401 & 1.404\\ 800.0 & 4.928 & 0.983 & 23168.333 & 159.55 & 29.623 & 8.837 & 97820.014 & 1.403\\ 900.0 & 5.536 & 0.994 & 26142.863 & 163.053 & 29.88 & 9.094 & 113121.903 & 1.401\\ 1000.0 & 6.151 & 1.006 & 29146.545 & 166.217 & 30.204 & 9.418 & 128756.479 & 1.399\\ 1100.0 & 6.774 & 1.019 & 32185.337 & 169.113 & 30.579 & 9.793 & 144693.583 & 1.397\\ 1200.0 & 7.406 & 1.034 & 35263.602 & 171.792 & 30.99 & 10.204 & 160909.037 & 1.395\\ 1300.0 & 8.051 & 1.051 & 38384.094 & 174.289 & 31.422 & 10.635 & 177383.01 & 1.392\\ 1400.0 & 8.709 & 1.069 & 41548.164 & 176.634 & 31.86 & 11.074 & 194098.888 & 1.389\\ 1500.0 & 9.382 & 1.089 & 44756.054 & 178.847 & 32.297 & 11.511 & 211042.494 & 1.386\\ 2000.0 & 13.015 & 1.193 & 61416.861 & 188.419 & 34.278 & 13.492 & 298792.538 & 1.371\\ 2500.0 & 17.181 & 1.299 & 78962.144 & 196.243 & 35.838 & 15.051 & 390859.016 & 1.357\\ 3000.0 & 21.964 & 1.397 & 97201.262 & 202.89 & 37.077 & 16.291 & 486526.149 & 1.345\\ 3500.0 & 27.428 & 1.486 & 116007.188 & 208.686 & 38.122 & 17.336 & 585293.51 & 1.335\\ 4000.0 & 33.632 & 1.568 & 135303.791 & 213.838 & 39.047 & 18.26 & 686790.847 & 1.326\\ 4500.0 & 40.634 & 1.644 & 155035.207 & 218.485 & 39.859 & 19.073 & 790733.265 & 1.318\\ 5000.0 & 48.493 & 1.713 & 175137.568 & 222.721 & 40.522 & 19.736 & 896893.167 & 1.311\\ 6000.0 & 66.995 & 1.831 & 216078.024 & 230.183 & 41.192 & 20.405 & 1.11513346e6 & 1.3\\ 7000.0 & 89.472 & 1.919 & 257182.729 & 236.52 & 40.847 & 20.061 & 1.340256372e6 & 1.292\\ 8000.0 & 116.05 & 1.972 & 297482.384 & 241.903 & 39.641 & 18.855 & 1.57122686e6 & 1.288\\ 9000.0 & 146.624 & 1.994 & 336293.379 & 246.476 & 37.931 & 17.145 & 1.807163794e6 & 1.286\\ 10000.0 & 180.9 & 1.99 & 373299.41 & 250.377 & 36.071 & 15.285 & 2.047326615e6 & 1.287\\ 11000.0 & 218.462 & 1.966 & 408479.598 & 253.731 & 34.302 & 13.516 & 2.29110708e6 & 1.288\\ 12000.0 & 258.842 & 1.93 & 442001.997 & 256.649 & 32.799 & 12.012 & 2.538015238e6 & 1.292\\ 13000.0 & 301.576 & 1.887 & 474137.426 & 259.222 & 31.549 & 10.763 & 2.787661831e6 & 1.295\\ 14000.0 & 346.243 & 1.84 & 505202.36 & 261.525 & 30.652 & 9.865 & 3.039740575e6 & 1.299\\ 15000.0 & 392.49 & 1.794 & 535523.749 & 263.617 & 30.079 & 9.293 & 3.294012259e6 & 1.304\\ 16000.0 & 440.045 & 1.75 & 565417.46 & 265.546 & 29.776 & 8.99 & 3.550291223e6 & 1.308\\ 17000.0 & 488.713 & 1.711 & 595174.689 & 267.35 & 29.787 & 9.0 & 3.808434146e6 & 1.311\\ 18000.0 & 538.38 & 1.676 & 625053.198 & 269.058 & 30.023 & 9.237 & 4.068330833e6 & 1.315\\ 19000.0 & 588.994 & 1.648 & 655271.777 & 270.692 & 30.442 & 9.656 & 4.329896651e6 & 1.318\\ 20000.0 & 640.561 & 1.625 & 686007.275 & 272.268 & 39.406 & 18.62 & 4.593066343e6 & 1.32\\ \end{tabular} \label{tab:result-thermo-Qpara} \end{table*}}
16
7
1607.04479
Context. Hydrogen is the most abundant molecule in the Universe. Its thermodynamic quantities dominate the physical conditions in molecular clouds, protoplanetary disks, etc. It is also of high interest in plasma physics. Therefore thermodynamic data for molecular hydrogen have to be as accurate as possible in a wide temperature range. <BR /> Aims: We here rigorously show the shortcomings of various simplifications that are used to calculate the total internal partition function. These shortcomings can lead to errors of up to 40 percent or more in the estimated partition function. These errors carry on to calculations of thermodynamic quantities. Therefore a more complicated approach has to be taken. <BR /> Methods: Seven possible simplifications of various complexity are described, together with advantages and disadvantages of direct summation of experimental values. These were compared to what we consider the most accurate and most complete treatment (case 8). Dunham coefficients were determined from experimental and theoretical energy levels of a number of electronically excited states of H<SUB>2</SUB>. Both equilibrium and normal hydrogen was taken into consideration. <BR /> Results: Various shortcomings in existing calculations are demonstrated, and the reasons for them are explained. New partition functions for equilibrium, normal, and ortho and para hydrogen are calculated and thermodynamic quantities are reported for the temperature range 1-20 000 K. Our results are compared to previous estimates in the literature. The calculations are not limited to the ground electronic state, but include all bound and quasi-bound levels of excited electronic states. Dunham coefficients of these states of H<SUB>2</SUB> are also reported. <BR /> Conclusions: For most of the relevant astrophysical cases it is strongly advised to avoid using simplifications, such as a harmonic oscillator and rigid rotor or ad hoc summation limits of the eigenstates to estimate accurate partition functions and to be particularly careful when using polynomial fits to the computed values. Reported internal partition functions and thermodynamic quantities in the present work are shown to be more accurate than previously available data. <P />The full datasets in 1 K temperature steps are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (<A href="http://130.79.128.5">http://130.79.128.5</A>) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/595/A130">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/595/A130</A>
false
[ "accurate partition functions", "internal partition functions", "New partition functions", "excited electronic states", "experimental values", "ad hoc summation limits", "thermodynamic quantities", "anonymous ftp", "molecular hydrogen", "thermodynamic data", "direct summation", "various simplifications", "protoplanetary disks", "the total internal partition function", "para hydrogen", "the estimated partition function", "polynomial fits", "previous estimates", "case", "rigid rotor" ]
11.542667
11.772592
180
1367063
[ "Greenbaum, Alexandra Z.", "Sivaramakrishnan, Anand" ]
2016OExpr..2415506G
[ "In-focus wavefront sensing using non-redundant mask-induced pupil diversity" ]
8
[ "-", "-" ]
[ "2016SPIE.9904E..0FP", "2018JATIS...4a9003D", "2018OptCo.425...85K", "2020Photo...8....3Q", "2022RemS...14.4681W", "2023ApPhL.123c4101Z", "2023PASP..135a5003S", "2023arXiv230904934D" ]
[ "astronomy", "physics" ]
5
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "1973JPhD....6L...6M", "1976JOSA...66..207N", "1982ApOpt..21.2758F", "1986JOSAA...3.1897F", "1995ApOpt..34.4951K", "1999JOSAA..16.1831F", "2004OptL...29.2840U", "2006SPIE.6265E..0RA", "2006SPIE.6265E..0XC", "2006SSRv..123..485G", "2008ApJ...688..701S", "2008SPIE.7010E..3CB", "2009SPIE.7440E..0YS", "2012OExpr..2029457C", "2012SPIE.8442E..2RD", "2012SPIE.8442E..3CA", "2013A&A...558A..33A", "2013PASP..125..422M", "2014MNRAS.440..125P", "2014OExpr..2212924C", "2014SPIE.9143E..3XP", "2015ApJ...798...68G", "2015JATIS...1b9001C" ]
[ "10.1364/OE.24.015506", "10.48550/arXiv.1607.01776" ]
1607
1607.01776_arXiv.txt
Intensity data from far-field or in-focus imagery of a point source is often used to determine the aberration in an optical system. Sometimes the aberration is an engineering or calibration quantity, as in the case of the immediately post-launch \textit{Hubble Space Telescope} (HST), and sometimes it is interesting of itself, as occurs in lensless X-ray microscopy. The Gerchberg-Saxton (GS) algorithm \cite{GS72} iteratively accomplishes an estimation of phase from image intensity and a knowledge of the pupil geometry, and does typically converge, but it can seem to converge to a spurious solution that is a local minimum, or converge to either of the two ambiguous solutions that are global minima. The unconstrained GS algorithm can converge to either the true pupil field $P(x)$ or its complex conjugate whose argument's sign is reversed, $P^*(-x)$ \cite{1986JOSAA...3.1897F}. Breaking the two-fold ambiguity can be accomplished in different ways. A consideration of more data taken with added defocus (or other kinds of phase diversity) can result in obtaining the true aberration \cite{Misell1973,1995ApOpt..34.4951K}. Extra data in the form of pupil diversity can also be used to break the sign ambiguity of the phase. Our approach is included amongst the latter class of methods. The GS phase retrieval method iteratively applies the known pupil transmission constraint in the pupil domain and the measured image intensity constraint in the image domain (e.g.~\cite{1981JOSA...71.1641F}). Stepping from one domain to the other is accomplished by a Fourier transform of pupil plane or image plane complex amplitudes. In the special case where the aberration is \textit{only} composed of functions where $P(x) = P^*(-x)$ (such as purely tip/tilt or coma) the phase can be recovered unambiguously with the unconstrained GS algorithm In practice the algorithm can converge to a local minimum rather than the true pupil phase unless the initial guess at the pupil phase is in some heuristic sense fairly close to the correct value. \begin{figure*}[htbp!] \centering \includegraphics[scale=0.425]{fig1a.pdf} \includegraphics[scale=0.3]{fig1b.pdf} \caption{JWST NIRISS pupil optics NRM and CLEARP. A life-sized prototype of the NRM is shown on the left, and the flight mask CLEARP in the NIRISS Pupil Wheel on the right. The JWST primary mirror is reimaged to a 40~mm diameter pupil in the plane of the NIRISS Pupil Wheel. The high quality of the image of the primary in NIRISS's entrance pupil \cite{2008SPIE.7010E..3CB}, makes NIRISS well-suited for our wavefront sensing technique during both commissioning and routine science operations.} \label{fig:pupil} \end{figure*} \begin{figure*}[htbp!] \centering \includegraphics[scale=0.65]{fig2.pdf} \caption{An example of true phase aberration, the quantity we are trying to measure, is shown at the top left. Exposures with the full pupil and the NRM produce aberrated PSFs. The NRM PSF is used to estimate the phases over each hole. We assume a known, fixed, pupil support $A(x)$, enforcing the pupil amplitude to match the known geometry of the pupil in each iteration. In addition, we also replace phases in the region shown with the average phase measured over each hole $\phi_{NRM}(x)$, measured from NRM fringe phases (bottom left two panels). We propagate between pupil and image planes since $P(x)\rightleftharpoons E_{\mathrm{image}}(k)$. In each iteration we replace the image field amplitude $a(k)$ at the focus by the square root of the PSF intensity (bottom right). We show an example of the estimated pupil phase at the 20th iteration. In this study we remove the NRM constraint after a fixed number of iterations (typically 100).} \label{fig:general} \end{figure*} We describe a phase retrieval method that uses a pair of in-focus images, one taken with the full pupil and the other with a non-redundant mask (NRM) in the pupil. A NRM consists of a set of holes in a mask where no hole-to-hole vector is repeated (e.g., Fig. \ref{fig:pupil} (left)). Our approach can be backup method for fine wavefront sensing that may typically be required in order to maintain diffraction-limited $2$~$\mu \mathrm{m}$ image quality on the 6.5~m 18-segment infrared James Webb Space Telescope (JWST), especially during routine operation and later stages of commissioning. A flight-ready wavefront sensing scheme using JWST's \textit{Near Infrared Camera} (NIRCam) has already been developed and tested \cite{2006SPIE.6265E..0RA}. However, JWST’s Near Infrared Imager and Slitless Spectrograph (NIRISS) \cite{2012SPIE.8442E..2RD} using its two pupil masks, NRM and CLEARP \cite{2009SPIE.7440E..0YS} (see Fig. 1) can serve as a backup sensing method. We show that these two pupil masks can be used to measure the telescope’s aberrations without introducing focus diversity by sweeping the secondary mirror through focus or placing some of NIRCam's three weak lenses in the beam. Our backup method reduces mission risk since it provides a second instrument that can measure JWST's wavefront. Using both NIRCam and NIRISS also provide wavefront measurements at different field points, which could assist with secondary mirror alignment when commissioning JWST. During routine astronomical observations a pair of images taken with the CLEARP and NRM pupil masks in one filter can provide a full wavefront measurement to interferometric accuracy. Such measurements could also support image deconvolution methods at all wavelengths in NIRISS, where all powered imaging optics are reflective. Using in-focus imagery and hardware optimized for science removes the need for dedicated wavefront sensing hardware, such as weak defocusing lenses, on future space telescopes. Our approach does not solve the persistent problem of non-common path wavefront aberrations in coronagraphs between the wavefront sensor and the focal plane mask \cite{2008ApJ...688..701S}, because we measure the wavefront at the science detector (rather than at the focal plane containing the coronagraphic occulter). We quantify the method's performance when faced with realistic limits of noise, size of the image (number of resolution elements), and wavefront error expected during certain commissioning phases of JWST. \S\ref{sec:motivation} motivates our approach and the design choices of our study. \S\ref{sec:mono}-\ref{sec:poly} show examples of how the algorithm performs under different conditions using both monochromatic and 8\%~bandwidth images, matching NIRISS's F480M filter on a continuous circular pupil. In \S\ref{sec:jwst} we apply the algorithm to a JWST-like pupil with segment tip, tilt and focus aberrations, as well as global pupil aberrations. We discuss our results in \S\ref{sec:discussion} in the context of JWST mirror phasing.
} Our constrained Gerchberg-Saxton approach to phase retrieval provides an efficient way to measure wavefront aberrations on current and future space telescopes using only in-focus images. The algorithm is suited to both continuous and segmented/obstructed pupils, though segment obstructions appear to limit the performance in our simple implementation. We use the first 15 Zernike or Hexike polynomial to smooth the wavefront in the pupil. We smooth over the full pupil in the case of a continuous pupil, and segment-wise in the case of the segmented pupil, using the appropriate set of polynomials. Alternative basis functions could handle different wavefront errors better. We ignore the edge effects of thin pupil obstructions. Mitigating the errors introduced by these effects may take more study. In this paper we have discussed our wavefront retrieval approach primarily in the context of commissioning JWST, although the method could provide wavefront knowledge during JWST science observations (for example, to support image deconvolution). On JWST this work may help avoid some secondary mirror focus sweeps during the commissioning phase of the telescope. We have presented a proof of concept of our method; further optimization of the method can be tailored to individual cases. On JWST, two in-focus exposures containing $10^6$ photons each (requiring exposure times of $<1$s full pupil images and seconds in NRM images for a star of $7.5^{\mathrm{th}}$ magnitude through the F480M filter) will provide enough signal to measure the wavefront errors of $\sim100$ nm, even in the presence of jitter, finite bandwidth and limited image size. Chromatic smearing, and finite image size are larger sources of error than the anticipated pointing jitter of JWST. Our method can tolerate up to 16 mas of pointing jitter, which is twice as large as JWST's required pointing accuracy. Frequent monitoring of mirror segment drift can be measured with our approach on NIRISS (which contains the 7-hole NRM used in this study) as a part of normal telescope operations. This provides a complimentary capability to trend wavefront stability over time in NIRISS alongside the main wavefront sensing monitoring program using NIRCam. We have focused on using a non-redundant mask to break the phase degeneracy in the unconstrained GS algorithm. It may also be possible to accomplish this with a redundant pupil that possesses asymmetries, using the asymmetric pupil wavefront sensor (APWFS) method \cite{Martinache_APWFS}, which also uses in-focus images. The APWFS algorithm could be used on data from other JWST instruments, such as NIRCam or MIRI. Our method has some key differences compared to the differential optical transfer function approach to wavefront sensing. The dOTF method similarly requires two in-focus images, one with a pupil modification, though the modification must be small. For JWST, Codona \cite{2015JATIS...1b9001C} suggests using small motions of the pupil wheel to block a portion of the pupil and achieve its required pupil diversity, which is not possible for MIRI's rachet-mechanism filter wheel. Our constrained GS approach uses two standard filter settings on NIRISS. With the possibility of using APWFS measurements, a single image could suffice for doing constrained GS wavefront sensing with MIRI. Using these methods for both monolithic and segmented future telescopes (such as the Wide Field Infrared Survey Telescope \cite{WFIRST} or the High Definition Space Telescope \cite{HDST}) can utilize science hardware for wavefront sensing, which has obvious benefits for weight, cost, complexity, and scope.
16
7
1607.01776
Wavefront estimation using in-focus image data is critical to many applications. This data is invariant to a sign flip with complex conjugation of the complex amplitude in the pupil, making for a non-unique solution. Information from an in-focus image taken through a non-redundant pupil mask (NRM) can break this ambiguity, enabling the true aberration to be determined. We demonstrate this by priming a full pupil Gerchberg-Saxton phase retrieval with NRM fringe phase information. We apply our method to measure simulated aberrations on the segmented James Webb Space Telescope (JWST) mirror using full pupil and NRM data from its Near Infrared Imager and Slitless Spectrograph (NIRISS).
false
[ "NRM fringe phase information", "full pupil", "NRM data", "NRM", "James Webb Space", "many applications", "a non-redundant pupil mask", "simulated aberrations", "Near Infrared Imager", "Slitless Spectrograph", "NIRISS", "complex conjugation", "a non-unique solution", "Information", "the segmented James Webb Space Telescope", "a full pupil Gerchberg-Saxton phase retrieval", "mirror", "JWST", "focus", "Gerchberg" ]
14.299858
10.591981
10
12473506
[ "Mantz, A. B.", "Allen, S. W.", "Morris, R. G." ]
2016MNRAS.462..681M
[ "Cosmology and astrophysics from relaxed galaxy clusters - V. Consistency with cold dark matter structure formation" ]
23
[ "Kavli Institute for Particle Astrophysics and Cosmology, Stanford University, 452 Lomita Mall, Stanford, CA 94305, USA; Department of Physics, Stanford University, 382 Via Pueblo Mall, Stanford, CA 94305, USA", "Kavli Institute for Particle Astrophysics and Cosmology, Stanford University, 452 Lomita Mall, Stanford, CA 94305, USA; Department of Physics, Stanford University, 382 Via Pueblo Mall, Stanford, CA 94305, USA; SLAC National Accelerator Laboratory, 2575 Sand Hill Road, Menlo Park, CA 94025, USA", "Kavli Institute for Particle Astrophysics and Cosmology, Stanford University, 452 Lomita Mall, Stanford, CA 94305, USA; SLAC National Accelerator Laboratory, 2575 Sand Hill Road, Menlo Park, CA 94025, USA" ]
[ "2017A&A...607A..81B", "2017MNRAS.469.1476S", "2018ApJ...861...71S", "2018JCAP...02..005H", "2018MNRAS.480.2898C", "2018MNRAS.481..361L", "2019A&A...621A..39E", "2019MNRAS.486.4001R", "2019SSRv..215...25P", "2019cxro.book...10A", "2020MNRAS.493.4783Y", "2020arXiv200904256D", "2022A&A...663A..17S", "2022A&A...665A..24P", "2022A&A...666A..41E", "2022A&A...668A..65T", "2022MNRAS.510..131M", "2022RAA....22l5015G", "2022arXiv221210232G", "2023MNRAS.521..790D", "2024ApJ...967...14C", "2024MNRAS.528.7274H", "2024arXiv240504577P" ]
[ "astronomy" ]
5
[ "galaxies: clusters: general", "dark matter", "X-rays: galaxies: clusters", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1969Ap......5...67E", "1993Natur.366..429W", "1997ApJ...490..493N", "1998ApJ...503..569E", "1998ApJ...504....1B", "2001MNRAS.321..372J", "2001MNRAS.321..559B", "2002ApJ...567..716R", "2002ApJ...568...52W", "2002ApJ...573....7E", "2003A&A...398..879E", "2003ApJ...597L...9Z", "2004ApJ...607..125T", "2004MNRAS.349.1039N", "2004MNRAS.353..457A", "2004MNRAS.353..624D", "2004MNRAS.355.1091K", "2005MNRAS.364..665D", "2006A&A...456...55Z", "2006ApJ...640..691V", "2006MNRAS.368..518V", "2007ApJ...655...98N", "2007MNRAS.379..209S", "2007MNRAS.381.1450N", "2008ApJ...688..709T", "2008JCAP...08..006M", "2008MNRAS.383..879A", "2008MNRAS.387..536G", "2008MNRAS.390L..64D", "2009A&A...501...61E", "2009ApJ...692.1060V", "2009ApJ...707..354Z", "2010ApJ...708..645R", "2010ApJ...722...55N", "2010MNRAS.402...21N", "2010MNRAS.406.1759M", "2011ARA&A..49..409A", "2011ASL.....4..204B", "2011ApJ...732...44S", "2011ApJ...736...52H", "2011ApJ...740..102K", "2012MNRAS.423.3018P", "2012MNRAS.427.1322L", "2013ApJ...763..147B", "2013ApJ...765...24N", "2013ApJ...769L..35O", "2013JCAP...07..008H", "2014A&A...571A..20P", "2014MNRAS.440.2077M", "2014MNRAS.441..378L", "2014MNRAS.441.3359D", "2015ApJ...806....4M", "2015ApJ...814..120D", "2015MNRAS.446.2205M", "2015MNRAS.449..199M", "2016A&A...586A..43V", "2016A&A...590A.126A", "2016A&A...594A..24P", "2016ApJ...821..116U", "2016MNRAS.456.4020M", "2016MNRAS.457.1522A", "2016MNRAS.457.4340K", "2017ApJ...840..104S" ]
[ "10.1093/mnras/stw1707", "10.48550/arXiv.1607.04686" ]
1607
1607.04686_arXiv.txt
\label{sec:intro} The CDM paradigm, within which the majority of matter in the Universe is weakly interacting, has enjoyed great success in explaining astrophysical and cosmological data. The inclusion of CDM in the now standard \LCDM{} cosmological model makes clear predictions for, among other things, the mass function of gravitationally collapsed structures (e.g.\ \citealt{Jenkins0005260, Evrard0110246, Tinker0803.2706}), the ratio of baryonic mass to CDM in galaxy clusters (e.g.\ \citealt{Eke9708070, Kay0407058, Borgani0906.4370}), and the distribution of mass within the structures that form hierarchically in the Universe (e.g.\ \citealt{Bullock9908159, Navarro0311231, Navarro9611107, Gao0711.0746}). Observations testing the predicted mass function and cluster baryon fraction have largely validated CDM (e.g.\ \citealt{White1993Natur.366..429W, Bahcall9803277, Reiprich0111285, Ettori0211335, Ettori0904.2740, Allen0405340, Allen0706.0033, Allen1103.4829, Vikhlinin0812.2720, Mantz0909.3098, Mantz1407.4516, Rozo0902.3702, Sehgal1010.1025, Benson1112.5435, Hasselfield1301.0816, Planck1303.5080, Planck1502.01597}). Observational tests of the distribution of mass within bound structures is in some ways more challenging, requiring spatially resolved measurements of the gravitational potential of objects whose mass is predominately dark. Nevertheless, an extensive body of work now exists, using primarily X-ray and gravitational lensing observations of galaxy clusters \citep{Vikhlinin0507092, Voigt0602373, Zhang0603275, Schmidt0610038, Mandelbaum0805.2552, Host0907.1097, Newman1209.1391, Okabe1302.2728, Du1510.08193, Merten1404.1376, Shan1502.00313, van-Uitert1506.03817}. In this work, we revisit the subject using \Chandra{} X-ray observations of a sample of massive, highly dynamically relaxed clusters. Although not representative of the cluster population at large, these systems are an ideal laboratory for deriving three-dimensional mass profiles based on observations of the intracluster medium (ICM) because departures from hydrostatic equilibrium and systematics due to projection are minimized. Both simulations and direct calibration using weak gravitational lensing indicate that the overall bias in \Chandra{} X-ray mass estimates for this sample is small ($\ltsim 10$ per cent; \citealt{Nagai0609247, Applegate1509.02162}). The specific features of the cluster mass distribution that we consider here are (1) the concentration parameter of cluster mass profiles, and its dependence on mass and redshift; and (2) the shape of mass profiles, in particular departures from the baseline model defined by \citet*{Navarro9611107} (hereafter NFW). The selection of the sample of massive, relaxed clusters employed here is detailed in the first paper in this series (\citealt{Mantz1502.06020}; hereafter Paper~\morphpaper{}). Papers~\cosmopaper{}, \profpaper{} and \calpaper{} \citep{Mantz1402.6212, Mantz1509.01322, Applegate1509.02162} employ \Chandra{} data for this sample to respectively constrain cosmological parameters, through gas-mass fraction measurements; scaling relations and thermodynamics of the ICM; and the average bias of the X-ray mass determinations. Section~\ref{sec:data} reviews our procedure for constraining cluster mass profiles from X-ray data, for which complete details are available in Papers~\cosmopaper{} and \profpaper{}, and introduces the specific mass models employed in this work. In Section~\ref{sec:results}, we review the specific predictions of CDM for cluster mass profile concentrations and shapes, present our results, and compare them to others in the literature. Section~\ref{sec:conclusions} summarizes our conclusions. We assume a concordance flat \LCDM{} cosmology with $h=0.7$ and $\Omegam=0.3$ throughout. Unless otherwise noted, quoted parameter uncertainties correspond to the maximum-likelihood (i.e.\ shortest) interval enclosing 68.3 per cent posterior probability, and best-fitting values are the posterior modes.
\label{sec:conclusions} We present constraints on mass profile models for massive, relaxed galaxy clusters based on an analysis of X-ray data from \Chandra{}. This analysis assumes hydrostatic equilibrium between the ICM and the gravitational potential, and includes data only at cluster radii where we are confident that both departures from equilibrium and systematics due to background modeling are minimal. Assuming the NFW mass profile model, the measured concentration--mass relation has a power-law slope with mass of $\kappa_m=-0.16\pm0.07$, consistent with CDM simulations of cosmological structure formation. The measured relation is consistent with being constant as a function of redshift, a feature that is not seen in simulations; however, this is plausibly the result of our selection of the most dynamically relaxed clusters at a given redshift. Simulations including gas physics, on which the same selection procedure is replicated, would be required to perform a completely fair comparison. We detect an intrinsic scatter of $\sigma_{\ln c}=0.16\pm0.03$ in the concentration--mass relation. Two clear high-concentration outliers from the mean relation also have the largest central cooling luminosities in the sample, suggesting a role for baryonic physics in the scatter. However, the remaining clusters do not support a simple trend between concentration and cooling luminosity at fixed mass. When fitting a GNFW mass profile, where the slope at small radii is a free parameter, we find an average value that is consistent with the NFW model, $\bar{\beta}=1.02\pm0.08$. However, there is significant cluster-to-cluster scatter ($\sigma_\beta=0.22\pm0.07$), even within the relaxed sample of clusters analyzed here. For the Einasto profile, we similarly measure a mean value, $\bar{\alpha}=0.29\pm0.04$, which is consistent with CDM predictions for clusters of the mass studied here. In this case, we measure a larger fractional intrinsic scatter, $\sigma_\alpha=0.12\pm0.04$. Overall, our results confirm that the mass distribution within the most massive halos is in agreement with CDM predictions. There are clear opportunities to improve this analysis in the future, both by finding additional highly relaxed clusters and obtaining deeper X-ray data (especially to better constrain 3-parameter models like the GNFW and Einasto models). Although it is beyond the scope of this work, X-ray data for dynamically relaxed clusters at intermediate radii can potentially be combined with weak lensing (at larger radii, $\gtsim r_{500}$), and strong lensing or stellar velocity dispersions (at small radii, $\ltsim 0.5r_{2500}$), in the vein of e.g.\ \citet{Newman1209.1391}. While these additional data sets have their own challenges in the form of projection effects and velocity anisotropy, a careful combination for an appropriately chosen cluster sample potentially provides the most complete possible view of the mass distribution within clusters.
16
7
1607.04686
This is the fifth in a series of papers studying the astrophysics and cosmology of massive, dynamically relaxed galaxy clusters. Our sample comprises 40 clusters identified as being dynamically relaxed and hot in Papers I and II of this series. Here we use constraints on cluster mass profiles from X-ray data to test some of the basic predictions of cosmological structure formation in the cold dark matter (CDM) paradigm. We present constraints on the concentration-mass relation for massive clusters, finding a power-law mass dependence with a slope of κ<SUB>m</SUB> = -0.16 ± 0.07, in agreement with CDM predictions. For this relaxed sample, the relation is consistent with a constant as a function of redshift (power-law slope with 1 + z of κ<SUB>ζ</SUB> = -0.17 ± 0.26), with an intrinsic scatter of σ<SUB>ln c</SUB> = 0.16 ± 0.03. We investigate the shape of cluster mass profiles over the radial range probed by the data (typically ∼50 kpc-1 Mpc), and test for departures from the simple Navarro-Frenk-White (NFW) form, for which the logarithmic slope of the density profile tends to -1 at small radii. Specifically, we consider as alternatives the generalized NFW (GNFW) and Einasto parametrizations. For the GNFW model, we find an average value of (minus) the logarithmic inner slope of β = 1.02 ± 0.08, with an intrinsic scatter of σ<SUB>β</SUB> = 0.22 ± 0.07, while in the Einasto case we constrain the average shape parameter to be α = 0.29 ± 0.04 with an intrinsic scatter of σ<SUB>α</SUB> = 0.12 ± 0.04. Our results are thus consistent with the simple NFW model on average, but we clearly detect the presence of intrinsic, cluster-to-cluster scatter about the average.
false
[ "-0.16 ±", "CDM predictions", "cluster mass profiles", "massive clusters", "cosmological structure formation", "SUB", "cluster", "σ", "small radii", "X-ray data", "β = 1.02 ±", "kpc-1 Mpc", "an intrinsic scatter", "NFW", "=", "CDM) paradigm", "Papers I", "the cold dark matter", "Einasto parametrizations", "0.12 ±" ]
12.407551
4.633191
-1
12462297
[ "Celletti, Alessandra", "Gales, Catalin" ]
2016FrASS...3...11C
[ "A study of the lunisolar secular resonance 2dot{ω}+dot{Ω}=0" ]
9
[ "University of Roma Tor Vergata, Department of Mathematics, Roma, Italy", "University of Iasi, Department of Mathematics, Iasi, Romania" ]
[ "2016MNRAS.460..802L", "2017IJNLM..90..147C", "2017MNRAS.464.4063R", "2019Chaos..29j1106G", "2022AdSpR..70..125C", "2022AdSpR..70.1234T", "2022CeMDA.134....6D", "2023AdSpR..72.2460L", "2023PhyD..45633946L" ]
[ "astronomy" ]
4
[ "space debris", "Lunisolar secular resonance", "Eccentricity growth", "Astrodynamics", "Celestial mechanics", "Astrophysics - Earth and Planetary Astrophysics", "70F15", "37N05", "34D08" ]
[ "1962AJ.....67..300K", "1962GeoJ....6..271C", "1979PhR....52..263C", "1980RSPSA.372..243H", "1989CeMDA..46..287L", "1997CeMDA..67...41F", "2001CeMDA..81...81B", "2008CeMDA.100..267R", "2015AdSpR..56..388C", "2015AdSpR..56.2626R", "2015MNRAS.449.3522R", "2015arXiv151202178C", "2016CeMDA.124..335D" ]
[ "10.3389/fspas.2016.00011", "10.48550/arXiv.1607.03767" ]
1607
1607.03767_arXiv.txt
Thousands of man-made objects, abandoned during space missions or remnants of operative satellites, orbit around the Earth at different altitudes. Their size varies from larger pieces, like old satellites or rocket stages, to dust-size particles given by fragmentation of satellites or even by collision events, like the impact between Kosmos 2251 and Iridium 33 in 2009, or the destruction of Fengyun-1C in 2007. The dynamics of space debris strongly differs according to the altitude from the Earth. To this end, one distinguishes 4 main regions as follows: \begin{itemize} \item [$(i)$] the LEO (Low Earth Orbit) region spans the altitude from 0 to 2\,000 km; here the objects feel, in order of importance, the gravitational attraction of our planet, the dissipation due to the atmospheric drag, the Earth's oblateness effect, the attraction of Moon and Sun, and the solar radiation pressure; \item [$(ii)$] the MEO (Medium Earth Orbit) region goes from 2\,000 to 30\,000 km of altitude; the forces felt by the debris are like in LEO, except that there is no atmospheric drag; \item [$(iii)$] the GEO (Geostationary orbit) region is located around the value of 42\,164.17 km from the Earth's center; geostationary objects move with an orbital period equal to the rotational period of the Earth; \item [$(iv)$] HEO (High Earth orbit) region, refers to the space region with altitude above the geosynchronous orbit. \end{itemize} In this work we are interested in a particular type of motion, which corresponds to a so-called secular resonance. In particular, we consider the orbital elements which are solutions of the relation \beq{res1} 2\dot{\omega}+\dot{\Omega}=0\ , \eeq where $\omega$ denotes the argument of perigee of the debris and $\Omega$ its longitude of the ascending node. A relation like \equ{res1}, involving quantities moving on long time-scales, is called a \sl secular resonance. \rm By considering the variations of $\omega$ and $\Omega$ as just due to the effect of the main spherical harmonics of the geopotential, one can show that equation \equ{res1} can be written just in terms of the inclination. As shown in \cite{HughesI}, there can be several secular resonances which depend on the inclination only. Among such resonances, \equ{res1} represents a very interesting case, since it has been shown that it affects the dynamics of objects in the MEO region (\cite{aR08}, \cite{Sanchez15}, \cite{RDGF2015}). Chaotic motions arise from the interaction and overlapping of nearby resonances (\cite{aR15}, \cite{DRADVR15}, \cite{RDADRV15}). In this paper we introduce three different models with increasing complexity, apt to study the resonance \equ{res1}. The simplest model is described by a one degree-of-freedom autonomous Hamiltonian, which is obtained by averaging over the fast angles and by neglecting the rates of variation of the lunar longitude of the ascending node. This model provides the essential features, like the location of stable equilibria with large as well as with small libration amplitude. The growth of the eccentricity can be easily explained by this integrable model. In the second model one does not average over the fast angles, but still retains the assumption that the longitude of the ascending node of the Moon is constant. Circulation and libration regions can be located, as well as the chaotic separatrix, although the dynamics is very complicated: overlapping of resonances, bifurcations and, as a consequence, the existence of equilibria at large eccentricities as well as at small eccentricities, variation of the amplitude of the resonance. The last model includes the variation of the lunar longitude of the ascending node and shows that large chaotic regions can appear, contributing to an irregular variation of the orbital elements.
\label{sec:conclusions} Lunisolar resonances might contribute to shape the dynamics of small bodies around the Earth (\cite{sB01}, \cite{DRADVR15}, \cite{aR15}). Among such resonances, that corresponding to $2\dot\omega+\dot\Omega=0$ is responsible for the growth in eccentricity. To explain this phenomenon, we compare three different models with increasing complexity, obtained averaging over fast angles (model a)), or just by neglecting the rate of variation of $\Omega_M$ (model b)), or rather including the variation of $\Omega_M$ (model c)). A comparison among these models provide us with the ingredients which lead to chaos and which provide an increase of the eccentricity. By comparing the results of models a)-b)-c), we infer that the dynamics around the stable equilibria at large values of the eccentricity is well represented by all models. On the contrary, for small values of the eccentricity the effect of the variation of the lunar longitude of the node plays a relevant role and, even if it occurs on long time scales, cannot be neglected for an accurate description of the dynamics. Finally, it is worth noticing that the growth in eccentricity provoked by the resonance $2\dot\omega+\dot\Omega=0$ can be used as an effective strategy to move space debris into non-operative or graveyard orbits. \appendix
16
7
1607.03767
The dynamics of small bodies around the Earth has gained a renewed interest, since the awareness of the problems that space debris can cause in the nearby future. A relevant role in space debris is played by lunisolar secular resonances, which might contribute to an increase of the orbital elements, typically of the eccentricity. We concentrate our attention on the lunisolar secular resonance described by the relation 2dot{ω}+dot{Ω}=0, where ω and Ω denote the argument of perigee and the longitude of the ascending node of the space debris. We introduce three different models with increasing complexity. We show that the growth in eccentricity, as observed in space debris located in the MEO region at the inclination about equal to 56°, can be explained as a natural effect of the secular resonance 2dot{ω}+dot{Ω}=0, while the chaotic variations of the orbital parameters are the result of interaction and overlapping of nearby resonances.
false
[ "nearby resonances", "space debris", "lunisolar secular resonances", "eccentricity", "the lunisolar secular resonance", "the space debris", "the secular resonance", "interaction", "overlapping", "perigee", "2dot{ω}+dot{Ω}=0", "increasing complexity", "the nearby future", "MEO", "the orbital parameters", "the orbital elements", "the ascending node", "a natural effect", "the chaotic variations", "the MEO region" ]
7.890325
15.557544
-1
12406109
[ "Vera, Matias", "Alonso, Sol", "Coldwell, Georgina" ]
2016A&A...595A..63V
[ "Effect of bars on the galaxy properties" ]
32
[ "ICATE, UNSJ-CONICET, CC49, 5400, San Juan, Argentina", "Facultad de Ciencias Exactas, Físicas y Naturales, UNSJ-CONICET, 5400, San Juan, Argentina", "Facultad de Ciencias Exactas, Físicas y Naturales, UNSJ-CONICET, 5400, San Juan, Argentina" ]
[ "2017ApJ...845...93K", "2017MNRAS.469.1054A", "2018A&A...618A.149A", "2019A&A...622A.132M", "2019A&A...632A..51C", "2019BAAA...61..169P", "2019MNRAS.489.4992D", "2020A&A...641A..77C", "2020A&A...644A..38D", "2020AJ....159..186D", "2020MNRAS.492.4697N", "2020MNRAS.499.1116F", "2020arXiv200309357D", "2021A&A...653A..71F", "2021A&A...654A.135D", "2021ApJ...915...23C", "2021MNRAS.507.4389G", "2021MNRAS.508.4459F", "2022ApJ...939...40L", "2022ApJ...940....1T", "2022ApJ...941...93O", "2022MNRAS.509.6155R", "2022MNRAS.510.5164C", "2023A&A...676A.140O", "2023ApJ...947...80B", "2023ApJ...949...91Z", "2023MNRAS.519.4801C", "2023MNRAS.521.1775G", "2024MNRAS.530.1171C", "2024MNRAS.530.5072M", "2024arXiv240411656S", "2024arXiv240505960G" ]
[ "astronomy" ]
11
[ "galaxies: spiral", "galaxies: formation", "galaxies: evolution", "Astrophysics - Astrophysics of Galaxies" ]
[ "1973ugcg.book.....N", "1979A&A....80..155L", "1979ApJ...233...67R", "1981A&A....96..164C", "1983IAUS..100..243A", "1984MNRAS.210P..19B", "1985ApJ...288..438E", "1985MNRAS.213..451W", "1986MNRAS.221P..41H", "1987ApJ...323...91D", "1989ApJ...342..677E", "1990A&A...236..333H", "1990ApJ...356..416P", "1990Natur.345..679S", "1991rc3..book.....D", "1992MNRAS.259..121V", "1992MNRAS.259..345A", "1993A&A...271..391C", "1993ApJ...411..137O", "1993RPPh...56..173S", "1994ApJ...420...87Z", "1994ApJ...424..599M", "1994ApJ...430L.105F", "1995A&A...301..649F", "1995AJ....109.2428M", "1996A&A...313...13H", "1996ASPC...91...54H", "1996ApJ...462..114N", "1996FCPh...17...95B", "1997A&A...323..363M", "1997ApJ...487..591H", "1998ApJ...493L...5D", "1999A&A...345...81C", "1999ApJ...510..125S", "1999ApJ...525..691S", "1999ApJ...527...54B", "2000A&A...356...89C", "2002A&A...392...83B", "2002AJ....124.1810S", "2002ARA&A..40..487F", "2003AJ....125.2348B", "2003ApJ...599L..29C", "2003MNRAS.341.1179A", "2003MNRAS.346.1055K", "2004A&A...415..941E", "2004ARA&A..42..603K", "2004ApJ...613..898T", "2004ApJ...615L.101B", "2004MNRAS.347..220B", "2004MNRAS.351.1151B", "2005ApJ...628..678D", "2005ApJ...630..837J", "2005ApJ...632..217S", "2005ApJ...634.1032K", "2005MNRAS.363..496A", "2006ApJ...637..214M", "2006ApJ...644..813E", "2006ApJ...645..209D", "2006ApJS..163..270G", "2007A&A...465L...9P", "2007ApJ...657..790M", "2007MNRAS.381..401L", "2007NCimB.122..935M", "2008ApJ...672L.107E", "2008ApJ...675.1141S", "2008ApJ...675.1194B", "2008ApJ...679L..73M", "2009A&A...494..891A", "2009A&A...495..775P", "2009A&A...501..207J", "2009ApJ...692L..34L", "2009ApJ...698.1639M", "2009ApJS..182..543A", "2009MNRAS.397..748P", "2010MNRAS.404..792M", "2010MNRAS.404.2087B", "2010MNRAS.405..783M", "2011MNRAS.410..166L", "2011MNRAS.411.2026M", "2011MNRAS.416.2182E", "2012ApJ...750..141L", "2012ApJS..198....4O", "2012MNRAS.420.1092D", "2012MNRAS.423.1485S", "2012MNRAS.424.2180M", "2013A&A...549A.141A", "2013MNRAS.431.2397D", "2013MNRAS.431.2560M", "2014A&A...570A...6S", "2014A&A...572A..86A", "2014MNRAS.437.1199C", "2015AJ....149....1Z", "2016MNRAS.457..917J" ]
[ "10.1051/0004-6361/201628750", "10.48550/arXiv.1607.08643" ]
1607
1607.08643_arXiv.txt
Galactic bars are structures observed in a significant fraction of spiral galaxies and are believed to have an important role in the dynamical evolution of their hosts. Several simulations show that bars can efficiently transport gas from the outer zones to the innermost central regions of the barred galaxies \citep{wei85,deb98,ath03}. By interaction with the edges of the bars, the gas clouds suffer shocks producing angular momentum losses and allowing a flow of material toward central kiloparcec scale \citep{SBF90}. Moreover, some works show that bars can be destroyed by a large central mass concentration (Roberts et al. 1979; Norman et al. 1996; Sellwood \& Moore 1999; Athanassoula et al. 2005). This finding indicates that currently non-barred disc galaxies possibly had a bar in the past (Kormendy \& Kennicutt 2004), and also, that bars may be recurrent in the galaxy life (Bournaud \& Combes 2002; Berentzen et al. 2004; Gadotti \& Souza 2006). So, in this context bars formed at different times, and with different conditions, might be present in the barred disc galaxies. Due to the high efficiency of gas inflow, galactic bars can alter several properties of disc galaxies on relatively short timescales. In this sence, the presence of bars can affect the star formation activity, stellar population, colors, modify the galactic structure \citep{atha83,buta96} and change the chemical composition \citep{comb93,martin95}, contributing to the evolution process of the host galaxies (Ellison et al. 2011, Zhou et al. 2015). In addition, the inflow processes have also been considered an efficient mechanism for trigger active galactic nuclei (AGN) \citep{comb93,cor03,alonso13,alonso14}, and to form bulges or pseudo-bulges (e.g., Combes \& Sanders 1981; Kormendy \& Kennicutt 2004; Debattista et al. 2005, 2006; Martinez-Valpuesta et al. 2006; M\'endez-Abreu et al. 2008; Aguerri et al. 2009). With respect to the relation between bars and host galaxy colors from statistical analysis, different studies show diverse results. Several observational works found that the bar fraction, $f_{bar}$, is higher in later-type spiral galaxies that are bluer and less concentrated systems (e.g. Barazza et al. 2008, Aguerri et al. 2009). However, other studies displayed an excess of barred galaxies with redder colors from different samples. \cite{master10a} found a high fraction of bars in passive red spiral galaxies for a sample obtained from the Galaxy Zoo catalog \citep{lintott11}. In addition, \cite{oh12} showed that a significant number of barred galaxies are redder than their counterparts of unbarred spiral galaxies. Recently, in our previous works (Alonso et al. 2013, 2014) we found an excess of red colors in spiral barred AGN with respect to unbarred active galaxies in a suitable control sample. The role of the bars on star formation and metallicity have been the subject of several works, showing unclear conclusions. Many studies found that bars enhanced the star formation rates (SFR) in spiral galaxies compared with unbarred ones (e.g. Hawarden et al. 1986; Devereux 1987; Hummel 1990), while other works show that bars do not guarantee increase in star formation activity (Pompea \& Rieke 1990; Martinet \& Friedli 1997; Chapelon et al 1999). In the similar way, different authors found diverse results in the metalicity studies in barred galaxies with respect to their unbarred counterparts (e.g. Vila$-$Costas \& Edmunds 1992; Oey \& Kennicutt 1993; Martin \& Roy 1994; Zaritsky et al 1994; Considere et al. 2000, Ellison et al. 2011). More recently, by using data from the CALIFA survey, S\'anchez-Bl\'azquez et al (2014) performed a comparative study of the stellar metallicity and age gradients in a sample of 62 spiral galaxies finding no differences with the presence or absence of bars. The discrepancy in the results of the bar effects on SFR and metallicity may depend on the host galaxy morphology (Huang et al. 1996; Ho et al. 1997; James et al. 2009) and may be also due to the length$/$strength of the bar (Elmegreen \& Elmegreen 1985, 1989; Erwin 2004; Menendez-Delmestre et al. 2007). Similarly, some studies (e.g. Athanassoula 1992; Friedli et al. 1994; Friedli \& Benz 1995) from numerical simulations found such trends, showing that bar strength is related to the efficiency and quantity of gas inflow, and therefore with the star formation activity and metallicity gradients. Furthermore, different authors have proposed diverse ways to build control samples from unbarred galaxies, used to obtain conclusions from comparative studies, and so the discrepancy in the results could be due to a biased selection of these samples. In this direction, \cite{perez09} found that a control sample for interacting galaxies should be selected matching, at least, redshift, morphology, stellar masses, and local density environment. This is also a suitable criteria for building control samples of barred galaxies (Alonso et al. 2013, 2014). Motivated by these finds, in this paper we conducted a detailed analysis of the effect of bars on host galaxy properties, with respect to the unbarred ones by studying different characteristics (e.g. color, stellar population, star formation activity, metallicity) with the aim of assessing whether bar structure in discs play a significant role in modifying galaxy properties, and how is this effect. This paper is structured as follows. Section 2 presents the procedure used to construct the catalog of barred galaxies from Sloan Digital Sky Survey (SDSS), the classification of the bar structures and the control sample selection criteria. In section 3, we explore different properties of the barred spirals, in comparison with unbarred galaxies obtained from a suitable control sample. We analize in details the influence of bars on star formation activity, stellar population, color indexes and metallicity in host spiral galaxies, with respect to unbarred ones. Finally, section 4 summarizes our main results and conclusions. The adopted cosmology through this paper is $\Omega = 0.3$, $\Omega_{\lambda} = 0.7$, and $H_0 = 100~ \kms \rm Mpc$.
\begin{enumerate} \item We found 522 strong-barred, 770 weak-barred, and 3711 non-barred galaxies, which represents a bar fraction of 25.82$\%$, with respect to the full sample of spiral galaxies. This fraction agrees with previous studies found by other authors by visual inspection of different galaxy samples from optical images (Nilson 1973, de Vaucouleurs et al. 1991, Marinova et al. 2009, Masters et al. 2010b, Alonso et al. 2013). \item We observed that strong-barred galaxies show lower star formation activity and older stellar populations, with respect to weak-barred and unbarred disc objects. We also found a significant fraction ($\approx20\%$) of strong-barred galaxies with older stellar population and low efficient star formation rate that have lenticular morphology (SB0 type). This result shows that, when S0 galaxies contain bar, it is usually a strong structure, in agreement with Aguerri et al. (2009). \item We also studied the star formation activity and the age of stellar populations of galaxies as a function of $log(M_{*})$ and concentration index, $C$, in barred galaxies with weak/strong bars, and in the control sample. We found that strong-barred galaxies show a systematically lesser efficient star formation activity and older stellar population for different stellar mass bins, and towards earlier morphology, with respect to the other samples of galaxies with weak and without bars. \item We examined the color distributions of different samples studied in this work, and we found that there is a significant excess of strong barred host galaxies with red colors. We also found that galaxies with strong bars are redder, for the whole concentration index range, with respect to their counterparts of weak-barred and unbarred disc objects. In particular, for strong barred galaxies that belong to the minor peaks of the star formation and stellar population distributions (see Fig. 3) these tendencies are clearly more significant, showing a high fraction of host galaxies with extremely red colors. These findings suggest that bar perturbations have a considerable effect in modifying galaxy colors in the host galaxies, producing an acceleration of the gas processing, when bar became prominent enough. \item The color-magnitude and color-color diagrams show that strong-barred galaxies are mostly grouped in the red region, while unbarred and weak-barred objects are more extended to the blue region. The positions in the color diagrams, could indicate the existence of an evolutive relation between the different considered galaxy type. In this scenario, an unbarred galaxy would begin to form a bar as a consequence of a gravitational disturbance in the disk. Then, matter would fall down to the center of the galaxy, making place to a weak bar which would become gradually more prominent while the inflow accumulates material in the center. At first, the weak bar would not be able to alter significantly the host characteristics, but then, when this structure is strong enough, it could affect significantly the galaxy properties. \item We also explore the metallicity, which principally reflects the amount of gas reprocessed by the stars. It shows that galaxies with strong bars present an important excess towards high metallicity values, while unbarred and weak-barred disc objects have similar distributions. The mass-metallicity relation reflects that although unbarred and weak-barred galaxies are fitted by similar curves, strong-barred ones show a curve which falls abruptly. It is more important in low stellar mass galaxies ($log(M_{*}/M_{\sun}) < 10.0$). This behaviour could be suggesting that prominent bars produce an accelerating effect on the gas processing, producing significant changes in the physical properties, also reflected in the evolutionary stages of the host galaxies. \end{enumerate}
16
7
1607.08643
<BR /> Aims: With the aim of assessing the effects of bars on disk galaxy properties, we present an analysis of different characteristics of spiral galaxies with strong bars, weak bars and without bars. <BR /> Methods: We identified barred galaxies from the Sloan Digital Sky Survey (SDSS). By visual inspection of SDSS images we classified the face-on spiral galaxies brighter than g&lt; 16.5 mag into strong-bar, weak-bar, and unbarred galaxies. With the goal of providing an appropriate quantification of the influence of bars on galaxy properties, we also constructed a suitable control sample of unbarred galaxies with similar redshifts, magnitudes, morphology, bulge sizes, and local density environment distributions to those of barred galaxies. <BR /> Results: We found 522 strong-barred and 770 weak-barred galaxies; this represents a bar fraction of 25.82% with respect to the full sample of spiral galaxies, in good agreement with several previous studies. We also found that strong-barred galaxies show lower efficiency in star formation activity and older stellar populations (as derived with the D<SUB>n</SUB>(4000) spectral index) with respect to weak-barred and unbarred spirals from the control sample. In addition, there is a significant excess of strong-barred galaxies with red colors. The color-color and color-magnitude diagrams show that unbarred and weak-barred galaxies are more extended towards the blue zone, while strong-barred disk objects are mostly grouped in the red region. Strong-barred galaxies present an important excess of high metallicity values compared to unbarred and weak-barred disk objects, which show similar distributions. Regarding the mass-metallicity relation, we found that weak-barred and unbarred galaxies are fitted by similar curves, while strong-barred ones show a curve that falls abruptly with more significance in the range of low stellar masses (log (M<SUB>∗</SUB>/M<SUB>⊙</SUB>) &lt; 10.0). These results would indicate that prominent bars produced an accelerating effect on the gas processing, reflected in the significant changes in the physical properties of the host galaxies.
false
[ "barred galaxies", "unbarred galaxies", "spiral galaxies", "galaxy properties", "disk galaxy properties", "weak bars", "strong bars", "Strong-barred galaxies", "strong-barred galaxies", "bars", "prominent bars", "unbarred and weak-barred galaxies", "similar distributions", "strong-barred disk objects", "similar curves", "low stellar masses", "several previous studies", "local density environment distributions", "older stellar populations", "strong-barred ones" ]
10.542663
6.851083
133
3815852
[ "Afshari, M.", "Peres, G.", "Jibben, P. R.", "Petralia, A.", "Reale, F.", "Weber, M." ]
2016AJ....152..107A
[ "X-Raying the Dark Side of Venus—Scatter from Venus’ Magnetotail?" ]
4
[ "Dipartimento di Fisica e Chimica, Università di Palermo, Piazza del Parlamento 1, I-90134, Italy; INAF—Osservatorio Astronomico di Palermo, Palermo, Piazza del Parlamento 1, I-90134, Italy;", "Dipartimento di Fisica e Chimica, Università di Palermo, Piazza del Parlamento 1, I-90134, Italy; INAF—Osservatorio Astronomico di Palermo, Palermo, Piazza del Parlamento 1, I-90134, Italy", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Dipartimento di Fisica e Chimica, Università di Palermo, Piazza del Parlamento 1, I-90134, Italy; INAF—Osservatorio Astronomico di Palermo, Palermo, Piazza del Parlamento 1, I-90134, Italy", "Dipartimento di Fisica e Chimica, Università di Palermo, Piazza del Parlamento 1, I-90134, Italy; INAF—Osservatorio Astronomico di Palermo, Palermo, Piazza del Parlamento 1, I-90134, Italy", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA" ]
[ "2017ApJ...844....3L", "2018ApJ...865..132D", "2019PASJ...71R...1H", "2023JCAP...03..001F" ]
[ "astronomy" ]
4
[ "planets and satellites: atmospheres", "planets and satellites: individual: Venus", "planets and satellites: terrestrial planets", "Sun: UV radiation", "Sun: X-rays", "gamma rays", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1972JOSA...62...55R", "1974AJ.....79..745L", "1997GeoRL..24.1163G", "1997dis..book.....J", "2000eaa..bookE3390F", "2001ApJ...556L..91S", "2002A&A...394.1119D", "2002PASP..114.1051S", "2007PASJ...59S.853W", "2007SoPh..243...63G", "2008P&SS...56.1414D", "2008SoPh..249..263K", "2009ApJ...690.1264D", "2011SoPh..269..169N", "2012SoPh..275....3P", "2012SoPh..275...17L", "2013ApJ...765..144P", "2014SoPh..289.2781K", "2015NatCo...6.7563R", "2016JSWSC...6A...1G", "2016SoPh..291..317T" ]
[ "10.3847/0004-6256/152/4/107", "10.48550/arXiv.1607.06697" ]
1607
1607.06697_arXiv.txt
Transits of Mercury and Venus across the solar disk are well-observed celestial phenomena. Recently, the transit of Mercury observed with \textit{Hinode}/X-Ray Telescope (XRT; \citealt{Gol07}) has been used by \citet{Web07} to test the sharpness of the instrument Point Spread Function (PSF). \citet{rea15} used \textit{Hinode}/XRT observations of a Venus transit to measure the size of Venus in the X-ray band thus inferring the extension and optical thickness of Venus\rq\ atmosphere. The methods and implications of the latter work reach into planetary physics and hint at similar methods to be potentially used, in the future, for exoplanets.\\ In this work we analyze the same set of observations to explore the residual X-ray emission observed in Venus\rq\ shadow and find, with the help of an updated version of the \textit{Hinode}/XRT PSF, that this emission is not due to instrumental scattering and may have an origin more directly related to Venus. Previous observations with Chandra in 2001 and then in 2006/2007 confirmed the X-ray emission from the sunlit side of the Venus (\citealt{Den02} and \citealt{den08}).\\ In Section 2 we present the observations of Mercury and Venus with a brief summary of the satellites and their instruments; in Section 3 we measure the residual flux in the shadow of Mercury in X-ray and of Venus in X-Ray, EUV and UV bands, and its evolution as Venus crosses the solar disk. In Section 4 we deconvolve X-ray images using the updated PSF and different codes, and again explore similarities and differences among the various observations; in Section 5 we describe the XRT straylight contamination and present our results taken with the Al-mesh filter. In Section 6 we show similar results obtained in EUV and UV bands. Section 7 contains our discussion and the conclusions.
We studied the Venus transit across the solar disk which occurred in 2012 and was observed with \textit{Hinode}/XRT in the X-Ray band and SDO/AIA in the EUV and UV bands. We have measured a significant X-Ray residual flux from Venus\rq\ dark side (i.e., from the Earth-facing side) during the transit that was significantly above the estimated noise level of 2 DN s$ ^{-1} $, as reported by \cite{Kob14}.\\ Let us discuss the systematic uncertainties of XRT flux. According to \cite{Kob14} there are two kinds of systematic uncertainties for XRT. The first are those which have a reliable quantitative correction procedure such as: dark current, Fourier, vignetting, and JPEG compression noise sources; their correction procedures have been successfully embedded in xrt\_prep.pro (the calibration reformatter). Since all of the data we have used for X-ray analysis (both Ti-poly and Al-mesh filters) have been prepared with xrt\_prep.pro, we expect that this class of uncertainties has been properly corrected and does not explain the observed residual.\\ The second kind of systematic uncertainties are model-dependent and are not included in xrt\_prep.pro; among them are light scattering by the grazing-incidence mirror of XRT, visible straylight leak and photon counting uncertainty \citep{Kob14}. In this respect the error bars for each flux value have been computed as follows. From the flux values, the exposure times, and the conversion factor from DN to photons, we have computed the number of photons collected and from them the statistical errors due to Poisson noise. This statistical error has been converted to a flux error bar per data point. The resulting error bar is typically less than the 2 DN s$ ^{-1} $ mentioned by \citep{Kob14}. In sections 4 and 5 we comprehensively discuss the PSF scattering and visible light leak effects.\\ To test the performance of the instrument's PSF (i.e., due to instrumental X-ray scattering) and the possible effect of the atmosphere on the residual flux, we studied a Mercury transit across the solar disk, observed with the \textit{Hinode}/XRT in 2006. We measured an apparent X-Ray residual flux in the case of Mercury before deconvolution.\\ For both Venus and Mercury we used a new version of the \textit{Hinode}/XRT PSF, selected well illuminated images in the X-Ray band, and deconvolved them. Even after deconvolution, flux from Venus\rq\ shadow has remained significant, while in the Mercury case it has become negligible. So it appears that the observed flux in Venus\rq\ shadow is real.\\ As for the Venus case, we have analyzed two X-ray datasets: a set collected with Ti-poly filter and another collected with Al-mesh filter. While the former is potentially strongly affected by a light leak that appeared a short time before that Venus transit, the latter is not. Both datasets, however, clearly show the presence of a significant flux from Venus\rq\ dark side, showing the reality of this effect. Although we consider the results from the Al-mesh data set to be a strong confirmation for the observed X-Ray residual flux, the Ti-poly results also provide more confidence about the observed residual flux and prove this effect in more than one filter. The level of the residual flux is not constant: as Venus crosses the solar disk it gradually grows, reaching a maximum roughly halfway through the transit, and then gradually decreases as it approaches the solar limb. The flux changes by an almost fixed factor of the flux of the surrounding solar regions (i.e., along nearby lines of sight) as shown in Figs.~\ref{LCTiDecon} and \ref{LCDeconAl}. On the other hand the use of the PSF and the test on Mercury convincingly shows the removal of any PSF effect. Furthermore, any light leak effect would instead be expected to be almost uniform or constant in time.\\ The PSF of XRT has also been determined at 1.0 keV. We find that deconvolving the images with this PSF reduces by a factor of about 0.5, on average, the flux inside the Venusian disk. More specifically, the mean flux across the disk before deconvolution is about 24 DN s$ ^{-1} $. After deconvolution with the 0.56 keV PSF model it is about 20 DN s$ ^{-1} $, but after deconvolution with the 1.0 keV PSF model it is about 10 DN s$ ^{-1} $. Therefore a significant flux level still remains, even after deconvolution with the 1.0 keV PSF, showing the reality of the effect nonetheless.\\ In this respect, however, we believe that the PSF at 0.5 keV is more appropriate to our study. In fact, we are detecting photons coming from the corona and re-processed at Venus or in Venus\rq\ magnetotail (or something related to Venus), a process which should not raise photon energy. The corona is at a few MK (at most 3 or 5, and only then in some places like active region cores), and no flare appears during the Venus transit. So a relatively smaller fraction of photons are expected even at 0.56 keV. We use the Al-mesh and Ti-poly filters; so considering the coronal spectrum folded with the Al-mesh filter response (Al-mesh data are the most reliable ones for Venus X-ray observations), we may shift the average of the observed plasma emission some 0.1 keVs closer to, but not at, 0.56 keV. On one hand we are confident that the general result is robust against the choice of PSF model, but the use of the 0.56 keV PSF can be considered to be a conservative evaluation, and so we can use the relevant results quite safely. \\ The analogous kind of analysis made in four EUV bands observed with SDO/AIA has shown that there is also some flux in these bands coming from Venus\rq\ night side, and that its evolution clearly follows that of the flux inside an annulus surrounding Venus. The light curves do not show, however, any trend similar to that of the X-Ray flux.\\ Past X-ray observations of Venus were very different, in many respects. In January 2001, Venus was observed for the first time with the Chandra X-ray telescope. \citet{Den02} proposed that the fluorescent scattering of solar X-rays from Venus\rq\ atmosphere was the primary source of the X-ray emission they observed. Not only the morphology, but also the observed X-ray luminosity was consistent with the scattering of solar X-rays \citep{Den02}. \\ In 2006 and 2007 again with Chandra, besides fluorescent scattering, Solar Wind Charge eXchange (SWCX) emission was clearly detected. Comparison of X-ray images taken in 2006 and 2007 with those obtained in 2001 (taken at a similar phase angle) showed that the limb brightening had increased. This would be the case if the X-ray radiation from Venus was the superposition of scattered solar X-rays and SWCX emission. The lack of detection of any SWCX-induced X-ray halo in the first Venus observation was explained by being during a high level of the solar X-ray cycle \citep{den08}.\\ Previous X-ray observations, however, have shown X-ray emission from the sunlit side of Venus. The low intensity we detect in X-ray and EUV comes from the dark side of Venus, and appears to have a totally different origin; it appears to evolve during the transit remaining, at any time, approximately proportional to the emission of the solar regions along nearby lines of sight. This intensity cannot be due to scattering in the upper atmosphere of Venus because we should detect a brighter inner rim in Venus\rq\ shadow.\\ The effect we are observing could be due to scattering or re-emission occurring in the shadow or wake of Venus. One possibility is due to the very long magnetotail of Venus, ablated by the solar wind and known to reach Earth's orbit \citep{Gru97}. This magnetotail could be side-illuminated from the surrounding regions and could scatter, or re-emit, the radiation; the cone of Venus shadow reaches up to $ 9.6 \times 10^5$ km away from Venus, leaving ample space ($\approx 4.5 \times 10^7 $ km) for side-illuminating the magnetotail. The emission we observe would be the reemitted radiation integrated along the magnetotail.\\ One wonders if such an effect is important for exoplanets, in particular for those Jupiter-size planets orbiting very close to their stars; they may have a very large ablated tail, especially if they do not have a magnetic field. To some extent, the study of these tails may help to understand, among other issues, the presence (or lack thereof) of magnetic fields.\\ Future work will study in more detail this phenomenon: we plan to study some faint structures present in the shadow and address possible physical mechanisms involved in generating the residual emission. \\
16
7
1607.06697
We analyze significant X-ray, EUV, and UV emission coming from the dark side of Venus observed with Hinode/XRT and Solar Dynamics Observatory/Atmospheric Imaging Assembly (SDO/AIA) during a transit across the solar disk that occurred in 2012. As a check we have analyzed an analogous Mercury transit that occurred in 2006. We have used the latest version of the Hinode/XRT point spread function to deconvolve Venus and Mercury X-ray images, to remove instrumental scattering. After deconvolution, the flux from Venus’ shadow remains significant while that of Mercury becomes negligible. Since stray light contamination affects the XRT Ti-poly filter data we use, we performed the same analysis with XRT Al-mesh filter data, not affected by the light leak. Even the latter data show residual flux. We have also found significant EUV (304 Å, 193 Å, 335 Å) and UV (1700 Å) flux in Venus’ shadow, measured with SDO/AIA. The EUV emission from Venus’ dark side is reduced, but still significant, when deconvolution is applied. The light curves of the average flux of the shadow in the X-ray, EUV, and UV bands appear different as Venus crosses the solar disk, but in any of them the flux is, at any time, approximately proportional to the average flux in a ring surrounding Venus, and therefore proportional to that of the solar regions around Venus’ obscuring disk line of sight. The proportionality factor depends on the band. This phenomenon has no clear origin; we suggest that it may be due to scatter occurring in the very long magnetotail of Venus.
false
[ "Venus", "disk line", "residual flux", "Venus and Mercury X-ray images", "significant EUV", "Venus’ shadow", "XRT", "UV bands", "instrumental scattering", "significant X", "sight", "Atmospheric Imaging Assembly", "dark side", "UV emission", "SDO", "Solar Dynamics Observatory", "-", "ray", "XRT Al-mesh filter data", "Mercury" ]
8.236312
15.143734
-1
12581239
[ "Klasen, Michael", "Lyonnet, Florian", "Queiroz, Farinaldo S." ]
2017EPJC...77..348K
[ "NLO+NLL collider bounds, Dirac fermion and scalar dark matter in the B-L model" ]
95
[ "Institut für Theoretische Physik, Westfälische Wilhelms-Universität Münster, Münster, Germany", "Southern Methodist University, Dallas, TX, USA", "Particle and Astroparticle Physics Division, Max-Planck-Institut für Kernphysik, Heidelberg, Germany" ]
[ "2016JHEP...12..081A", "2016JHEP...12..106A", "2016PhLB..763..269Q", "2016PhRvD..94g6008L", "2016PhRvD..94i5024H", "2016arXiv161005237M", "2016arXiv161006587L", "2017EPJC...77..889W", "2017JCAP...01..042A", "2017JCAP...04..016A", "2017JCAP...07..016C", "2017JCAP...11..020A", "2017JHEP...04..081A", "2017JHEP...04..164A", "2017JHEP...08..092C", "2017JHEP...10..169D", "2017PhLB..771..508A", "2017PhLB..772..825A", "2017PhLB..775..196A", "2017PhRvD..95c5025O", "2017PhRvD..95e5019A", "2017PhRvD..96i5032O", "2017PhRvD..96k5014N", "2017Prama..89...52B", "2017arXiv170803955S", "2017arXiv171107634B", "2018ChPhC..42j3101D", "2018EPJC...78..189N", "2018EPJC...78..203A", "2018EPJC...78..696D", "2018EPJC...78..839H", "2018EPJP..133..477S", "2018FrP.....6...40C", "2018JCAP...11..026S", "2018JHEP...03..122D", "2018JHEP...08..190E", "2018MPLA...3350034L", "2018MPLA...3350226M", "2018PhLB..781..561N", "2018PhLB..784..385C", "2018PhR...731....1L", "2018PhRvD..97a5001B", "2018PhRvD..97d3009A", "2018PhRvD..98c5027B", "2018PhRvD..98e5014F", "2018PhRvD..98i5019A", "2018ScPP....5...36B", "2018arXiv180105594C", "2018arXiv180306793O", "2018arXiv180803352B", "2018arXiv180902453E", "2018arXiv181012920B", "2019EPJC...79..574C", "2019JHEP...02..059B", "2019JHEP...05..154A", "2019JHEP...11..129F", "2019JHEP...12..109B", "2019PhRvD..99a5038B", "2019PhRvD.100a5042A", "2019PhRvD.100c5022O", "2019PhRvD.100e5027B", "2019PhRvD.100k5023D", "2019ScPP....6...20E", "2019arXiv190302572C", "2019arXiv190711556G", "2019arXiv190711557G", "2019arXiv191102419F", "2020EPJC...80..557N", "2020JHEP...04..049F", "2020NuPhB.95014841B", "2020PhLB..80635499G", "2020PhLB..80735537L", "2020PhRvD.101a5006G", "2020PhRvD.102i5021F", "2020arXiv200810627B", "2020arXiv201214855M", "2021JHEP...02..223F", "2021JHEP...05..150B", "2021JHEP...07..166B", "2021JPhG...48b5002C", "2021PhLB..82136605G", "2021PhRvD.103k5026A", "2022JHEP...05..182B", "2022NuPhB.98516028C", "2022PhRvD.105a5015B", "2022PhRvD.105c5016A", "2022arXiv220808462C", "2023JHEP...02..103B", "2023JHEP...03..182A", "2023NuPhB.98616057C", "2023PhRvD.107f3005H", "2023PhRvD.108g5021L", "2024JHEP...01..013D", "2024PhLB..84838382A", "2024arXiv240415987B" ]
[ "astronomy", "physics" ]
6
[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Experiment" ]
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[ "10.1140/epjc/s10052-017-4904-8", "10.48550/arXiv.1607.06468" ]
1607
1607.06468_arXiv.txt
\label{sec:1} The availability of data from collider, direct and indirect searches for dark matter has raised the importance of dark matter complementarity across these search strategies. In this context, effective field theories and simplified models have become popular tools, as they can capture most of the dark matter phenomenology. Planck measurements of the power spectrum of the cosmic microwave background radiation infer that the cold dark matter abundance should be around 27\% ($\Omega_{DM} h^2=0.12$), where $h$ is a parameter that accounts for uncertainties in the Hubble rate \cite{Ade:2013zuv}. This alone strongly constrains the viable parameter space of dark matter models. The observation of cosmic rays and gamma rays also offers a compelling probe for dark matter \cite{Hooper:2012sr,Bringmann:2012ez,Calore:2013yia,Kopp:2013eka,Galli:2013dna,Gomez-Vargas:2013bea,Berlin:2013dva,Madhavacheril:2013cna, Abazajian:2014fta,Bringmann:2014lpa,Gonzalez-Morales:2014eaa,Buckley:2015doa,Huang:2016pxg}. In particular, the Fermi-LAT sensitivity to continuous gamma-ray emission from dark matter annihilations taking place in Dwarf Galaxies resulted in restrictive bounds in the annihilation cross section today, namely $\sigma v < 3 \times 10^{-26} {\rm cm^3/s}$ for masses of 80 GeV and annihilation into $b\bar{b}$ quark pairs~\cite{Ackermann:2015zua}. This rules out a multitude of light WIMP (weakly interacting massive particles) models in which velocity-independent interactions occur. Moreover, underground detectors using liquid XENON, such as XENON \cite{Aprile:2012nq} and LUX \cite{Akerib:2013tjd} that use scintillation and ionization measurements to discriminate signal from background events, observed no excess, leading to the exclusion of spin-independent WIMP-nucleon scattering cross sections larger than $10^{-45} {\rm cm^2}$ for WIMP masses of $50$~GeV. Other experiments have placed complementary limits in particular at lower masses such as SUPERCDMS, which uses Ge targets \cite{Agnese:2015nto}. The ongoing XENON1T \cite{Aprile:2015uzo} and LZ \cite{Malling:2011va} experiments are expected to bring down the limits by roughly two orders of magnitude in the absence of signal and zero background events. Besides the indirect and direct detection probes, the Tevatron \cite{Bai:2010hh} and the LHC \cite{Goodman:2010ku,Fox:2011pm} have proven to be great laboratories to test dark matter models. In the case where the dark and visible sectors are connected by vector mediators, dijet \cite{An:2012va,Frandsen:2012rk,Alves:2013tqa,Fairbairn:2016iuf} and dilepton \cite{Profumo:2013sca,Alves:2015pea,Alves:2015mua,Allanach:2015gkd,Kahlhoefer:2015bea} bounds are by far the most stringent constraints. Dark matter phenomenology is then dictated by gauge interactions which are determined, once the gauge group behind the origin of the vector mediator is known. The common approach is to consider simplified lagrangians that encompass both Dirac and Majorana dark matter fermions and then to compute dark matter observables; namely, relic density, annihilation and scattering cross sections, the latter being spin-independent and spin-dependent for Dirac and Majorana fermions, respectively\footnote{Dirac fermions also induce spin-dependent interactions but the spin-independent ones lead to stronger constraints.}. The simplified dark matter model approach is interesting, intuitive and serves as a guide for future work. However, they might lead to different results once embedded in a complete theory. In the context of the B-L model, dark matter scenarios have been previously investigated. In~\cite{Li:2010rb}, the authors discussed the radiative see-saw mechanism to account for neutrino masses and focused exclusively on dark matter abundance. Supersymmetric B-L extensions~\cite{Khalil:2008ps,Burell:2011wh,Basso:2012gz} and a conformal approach~\cite{Okada:2012sg} have also been investigated. Even though later disfavored in~\cite{Mambrini:2015sia}, a global B-L symmetry has been proposed~\cite{Baek:2013fsa}. In~\cite{El-Zant:2013nta} a warm dark matter scenario was investigated. The possibility of having one of the right-handed neutrinos to be the dark matter candidate was entertained in \cite{Sahu:2005fe,Basak:2013cga,Kaneta:2016vkq}, whereas in \cite{Rodejohann:2015lca} an additional scalar played this role. This extra scalar dark matter was also investigated in \cite{Guo:2015lxa}, but in the context of classical scale invariance. The authors of \cite{Sanchez-Vega:2015qva,Patra:2016ofq} considered an exotic B-L model and advocated the presence of many scalar fields. Finally, the authors of \cite{Duerr:2015wfa} studied Dirac fermion dark matter in the context of a $U(1)_{B-L}$ symmetry, but with the inclusion of LEP bounds only they discussed gamma-ray lines emissions, which turned out to be irrelevant unless one lives very close to the resonance with a dark matter quantum number under B-L larger than three. Thus, our work supplements previous studies for the following reasons: (i) Both fermionic and scalar dark matter realizations are discussed as well as several quantum numbers and gauge couplings options. (ii) We investigate two-component dark matter scenarios. (iii) We perform a detailed collider study at next-to-leading order (NLO) plus next-to-leading logarithmic (NLL) accuracy using recent dilepton data from the LHC at 13 TeV, which are often ignored due to the handy LEP limits. (iv) Finally, the region of parameter space allowed/excluded by limits from the LHC, LEP and indirect detection experiments in dependence of the mass of the mediator, gauge couplings and dark matter mass is presented.
Supplementing the SM with an extra $\mathrm{U}(1)_{B-L}$ gauge symmetry is an appealing possibility. In this paper, we studied the dark matter phenomenology of simplified models exhibiting such a gauge symmetry and in particular the possibilities of having Dirac fermion as well as scalar dark matter with and without broken B-L symmetry. In this context, we determined the impact of constraints coming from indirect and direct detection experiments as well as collider limits. Bounds from LUX2015, LUX2016 and projected bounds from XENON1T have been considered along with the famous LEP limit. In addition, we re-interpreted dilepton searches from the LHC at 13 TeV and extracted competitive limits for the model. While XENON1T projected bounds have a very good potential to exclude most of if not all the parameter space for scalar dark matter, we found that Dirac fermion dark matter would still be viable in a larger region of the parameter space. Interestingly, it was shown that the LHC limits that were extracted from dilepton production are already better than the LEP bounds for small gauge couplings. Finally, we also considered a mixed dark matter scenario, in which the relic abundance is realized as a combination of both fermion and scalar dark matter. In this case, numerous points satisfying the required relic density, collider, direct and indirect dark matter constraints were found, showing that a minimal and successful two component dark matter model is realized.
16
7
1607.06468
Baryon and lepton numbers being accidental global symmetries of the Standard Model (SM), it is natural to promote them to local symmetries. However, to preserve anomaly-freedom, only combinations of B-L are viable. In this spirit, we investigate possible dark matter realizations in the context of the U(1)_B{-L} model: (i) Dirac fermion with unbroken B-L; (ii) Dirac fermion with broken B-L; (iii) scalar dark matter; (iv) two-component dark matter. We compute the relic abundance, direct and indirect detection observables and confront them with recent results from Planck, LUX-2016, and Fermi-LAT and prospects from XENON1T. In addition to the well-known LEP bound M_{Z^' }}/g_BL ≳ 7 TeV, we include often ignored LHC bounds using 13 TeV dilepton (dimuon + dielectron) data at next-to-leading order plus next-to-leading logarithmic accuracy. We show that, for gauge couplings smaller than 0.4, the LHC gives rise to the strongest collider limit. In particular, we find M_{Z^' }}/g_BL &gt; 8.7 TeV for g_BL=0.3. We conclude that the NLO+NLL corrections improve the dilepton bounds on the Z^' } mass and that both dark matter candidates are only viable in the Z^' } resonance region, with the parameter space for scalar dark matter being fully probed by XENON1T. Lastly, we show that one can successfully have a minimal two-component dark matter model.
false
[ "scalar dark matter", "possible dark matter realizations", "local symmetries", "accidental global symmetries", "both dark matter candidates", "SM", "XENON1T.", "Z^", "unbroken B-L", "broken B-L", "recent results", "a minimal two-component dark matter model", "prospects", "TeV", "LUX-2016", "(i) Dirac fermion", "(ii) Dirac fermion", "M_{Z^", "Planck", "the Standard Model" ]
8.513637
-1.757484
54
12311754
[ "López Fuentes, Marcelo", "Klimchuk, James A." ]
2015ApJ...799..128L
[ "Two-dimensional Cellular Automaton Model for the Evolution of Active Region Coronal Plasmas" ]
21
[ "Instituto de Astronomía y Física del Espacio, CONICET-UBA, CC. 67, Suc. 28, 1428 Buenos Aires, Argentina; Member of the Carrera del Investigador Científico, Consejo Nacional de Investigaciones Científicas y Técnicas (CONICET), Argentina.;", "NASA Goddard Space Flight Center, Code 671, Greenbelt, MD 20771, USA" ]
[ "2015BAAA...57..231L", "2015RSPTA.37340256K", "2016ApJ...817....5H", "2016ApJ...817L..11T", "2016ApJ...821...63B", "2016ApJ...828...86L", "2017ApJ...840....4A", "2017ApJ...842..108V", "2017ApJ...849...46V", "2018ApJ...853...82K", "2018BAAA...60..207L", "2020ApJ...899..156K", "2020ApJ...905..115S", "2020Chaos..30d3124L", "2020SoPh..295...21K", "2020SoPh..295..171N", "2021ApJ...906...59H", "2021FrASS...8...33D", "2022ApJ...936..113H", "2022ApJ...939...17L", "2024ApJ...965...53K" ]
[ "astronomy" ]
4
[ "Sun: activity", "Sun: corona", "Sun: magnetic fields", "Sun: X-rays", "gamma rays", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1978ApJ...220..643R", "1979ApJ...233..987V", "1983ApJ...264..635P", "1983ApJ...264..642P", "1988ApJ...330..474P", "1990PhyA..163..403B", "1991ApJ...380L..89L", "1991SoPh..133..357H", "1994ApJ...422..381C", "1995ApJ...446L.109L", "1998ApJ...501L.213K", "1999SoPh..187..229H", "2000ApJ...529..554P", "2000ApJ...530..999M", "2000ApJ...542.1088L", "2001ApJ...550.1036A", "2001SoPh..203..321C", "2002ApJ...572.1048A", "2002ApJ...576..533P", "2003ApJ...587..439W", "2004ApJ...605..911C", "2004IAUS..219..461B", "2005ApJ...622.1191D", "2006SoPh..234...41K", "2007A&A...462..311P", "2007ApJ...657.1127L", "2007SoPh..243....3K", "2007SoPh..243...63G", "2008A&A...492L..13B", "2008ApJ...682..654M", "2008ApJ...682.1351K", "2009ASPC..415..221K", "2009ApJ...704.1059D", "2010ApJ...711..228W", "2010ApJ...719..591L", "2010LRSP....7....5R", "2011ApJ...736..111T", "2011ApJ...738...24V", "2011soca.book.....A", "2012ApJ...752..161C", "2012ApJ...753...35V", "2012ApJ...756..126S", "2012ApJ...758...53B", "2012ApJ...759..141W", "2012RSPTA.370.3217P", "2012SoPh..275....3P", "2012SoPh..275...17L", "2013ApJ...771..115V", "2013ApJ...774...31G", "2014ApJ...783...12U", "2014ApJ...784...49C", "2014ApJ...795...76S" ]
[ "10.1088/0004-637X/799/2/128", "10.48550/arXiv.1607.03883" ]
1607
1607.03883_arXiv.txt
\label{intro} One of the most persistent conundrums of Solar Physics has been, and still is, the problem of coronal heating. The difficulties arise from both the theoretical and observational sides. Observationally, it is very difficult to explain with a single scenario the diverse set of observations obtained in different wavelengths. The first X-ray observations in the decade of 1970 suggested that coronal loops were in quasi-static equilibrium and that a steady or quasi-steady heating process was balanced by radiative and conductive losses (Rosner et al. 1978). However, the situation changed as soon as ultraviolet observations from space became available, especially when instruments like the Transition Region and Coronal Explorer (TRACE, Handy et al. 1999) began to produce higher resolution and cadence data. The evolution of many EUV loops was too dynamic and intermittent to be consistent with quasi-static evolutions. Furthermore, density determinations using EUV instruments showed that loops are too dense given their temperature and length and had scales heights that are too large to be quasi-static (Aschwanden et al. 2001, Winebarger et al. 2003). It was proposed then that EUV loops could be heated by impulsive mechanisms. Thermal evolution models based in this premise were successful in explaining many of the observed physical conditions and evolutions (see e.g., Klimchuk 2006, 2009; Reale 2010, and references therein). More recently it has been shown that impulsive heating plays an important role in the diffuse component of the corona as well (Viall \& Klimchuk 2011, 2012, 2013; Bradshaw et al. 2012, Warren et al. 2012, Schmelz \& Pathak 2012). The question is now if it is possible to understand X-ray loops, EUV loops, and the diffuse emission as part of the same phenomenon or if they are caused by completely different mechanisms. The problem has of course many different aspects: the emission evolution, geometry, location within active regions, physical conditions of the plasma, etcetera, and all of these have to be considered to provide an explanation. One proposed possibility is that all of the corona is heated by impulsive short duration events, but the rate of repetition is high at some locations (e.g., X-ray loops) and indiscernible from a continuous source (Warren et al. 2010). From the theoretical side, and in particular regarding the heating mechanism itself, several models have been proposed. They can be grossly classified in two types: models based on the dissipation of waves and models based on the dissipation of magnetic stresses (see the reviews by Klimchuk 2006, Reale 2010, Parnell \& De Moortel 2012). Both types can produce impulsive heating, but the best known idea is from the second category and came from Parker (1988). He proposed that loops are made of elementary magnetic strands that are shuffled and tangled by photospheric motions. Current sheets form at the boundaries between the strands, and energy is released by magnetic reconnection. Parker estimated that the energy content of a single impulsive event is roughly one-billionth of a large flare, so he coined the term ``nanoflare." We now use this term generically to describe any small spatial scale impulsive event, irrespective of the physical cause. Parker's mechanism, the basis of our investigation, could provide the impulsive events invoked in the previous paragraph to explain EUV loops, X-ray loops, and diffuse emission within the same phenomenological framework. Several studies over the years analyzed different aspects of the nanoflare heating problem, such as the conditions for reconnection (Parker 1983a, 1983b; Priest et al. 2002, Darlburg et al. 2005), the thermal evolution of the plasma (Cargill 1994, Cargill \& Klimchuk 2004) and the observed coronal emission (Warren et al. 2010, L\'opez Fuentes et al. 2007). Other relevant issues are the mechanism by which footpoint motions translate into magnetic stress, the geometrical and temporal characteristics of the interaction between strands and the energy distribution of the produced nanoflares. A series of models based on the concept of Self Organized Criticality (SOC, Bak et al. 1988) have been developed during the last 20 years to analyze this part of the problem (see e.g., Lu \& Hamilton 1991, Lu 1995, Longcope \& Noonan 2000, Morales \& Charbonneau 2008). The idea of this kind of approach is to use simple sets of rules for the injection of energy simulating the effect of footpoint motions, the interaction between magnetic strands, and the existence of a magnetic stress threshold beyond which energy release occurs (pedagogical reviews on the subject can be found in Charbonneau et al. 2001, and the book by Aschwanden 2011). In a recent paper we presented a simple pseudo-1D cellular automaton model to explain the observed evolution of soft X-ray loops (Lopez Fuentes \& Klimchuk 2010). Here we develop a more sophisticated model based on a similar aproach. Simulating the motions of magnetic strand footpoints in a 2D array we establish a series of rules for the interaction between strands and critical conditions for the magnetic stress relaxation through reconnection. The consequent energy release takes the form of short duration events whose output heats the plasma in the strands. To simulate the plasma response to these events we use the EBTEL code (Enthalpy Based Thermal Evolution of Loops, Klimchuk et al. 2008, Cargill et al. 2012). We analyze the output of the model by studying its dependence on the model parameters, the presence of power-laws in the energy distribution and the temporal properties of the produced nanoflares. In a second paper we will compare the results of the model with observed X-ray and EUV loops. The paper is organized as follows. In Section~\ref{model} we provide a detailed description of the model, the implementation of EBTEL to compute the plasma response and the obtainment of synthetic observations. In Section~\ref{analysis} we present and analyze the results and discuss some of their implications and conclude in Section~\ref{conclusions}.
\label{conclusions} We study the problem of nanoflare heating of coronal loops using a 2D cellular automaton (CA) model based on Parker's (1988) idea of footpoint shuffling and tangling of elementary magnetic strands. To determine the plasma response to the heating we use the Enthalpy Based Thermal Evolution of Loops (EBTEL) model. From the computed temperature and density evolutions and the known response of coronal observing instruments we simulate observed lightcurves. We study the dependence of the model's output on the relevant physical parameters and we find and analyze a series of predicted scalings that can be compared with future observations. Two primary results of our study concern the number distribution of nanoflares as a function of energy and the frequency with which nanoflares repeat on a given strand. We find that the number distribution obeys a power law with a slope of approximately -2.5. This is a robust result, with a standard deviation of 17\% as we vary the model parameters. For many years researchers have extrapolated power laws measured for flares and other resolvable events to lower energies in order to determine whether the corona could be heated by nanoflares. As pointed out by Hudson (1991), the slope of the distribution must be steeper than -2 in order for nanoflares to be energetically important. The results have been mixed, ranging from less steep (Aschwanden \& Parnell 2002 and references therein) to more steep (Krucker \& Benz 1998; Parnell and Jupp 2000; Benz 2004; Pauluhn \& Solanki 2007; Bazarghan et al. 2008), with shallower slopes tending to come from studies of flares and steeper slopes tending to come from studies of smaller impulsive events. The variation reflects the difficulty in measuring the slope, in part because of the assumptions that must be made in estimating the total energy that is released. We would also point out that there is no compelling reason to believe that the slope should be constant over the full range extending from large flares to nanoflares. From an observational perspective, the nanoflare repetition frequency is important only insofar as the delay between successive events is longer than, shorter than, or comparable to the plasma cooling time. In our models, the nanoflare frequency increases with loop length. Since hot (X-ray) loops are best explained by high-frequency nanoflares, and warm (EUV) loops are best explained by low-frequency nanoflares, this would suggest that hot loops should be longer than warm loops. This is not observed to be the case, however. Hot loops are more prevalent in the cores of active regions, while warm loops are more prevalent outside the core. Most core emission is contained in a diffuse component, however, rather than in observationally distinct loops (Viall \& Klimchuk 2011). Whether this diffuse emission is better explained by high or low frequency nanoflares is a matter of debate. See the studies summarized in Table 3 of Bradshaw et al. (2012) and the discussion of uncertainties in Guennou et al. (2013). Subramanian et al. (2014) have examined diffuse emission outside the core, and again the results are inconclusive regarding nanoflare frequency. Cargill (2014) recently suggested that all the results might be reconciled if nanoflares have random frequencies centered about a mean value that is comparable to a cooling time, i.e., intermediate frequency, and if there is a relationship between the event size and delay between events. The predictions of our model of course depend upon the assumptions that go into it. We have three main assumptions. First, magnetic reconnection does not occur unless a critical condition corresponding the misalignment angle between adjacent magnetic strands is met. This is supported by theoretical studies of the secondary instability of current sheets (Dahlburg et al. 2005, 2009). Second, the test for criticality does not happen until after a full stressing step has been made. In other words, reconnection holds off at least until photospheric convection has transported the magnetic footpoint a characteristic distance $d = 1$ Mm. The critical condition can therefore be exceeded, though not by a large amount. Third, reconnection, once it occurs, releases an amount of stress equal to the stress added during a driving step. It does not cause the total stress to drop far below the critical value. A strand can reconnect with multiple strands during a single iteration, but in practice the number rarely exceeds three. The mean is close to one. As a result the stress tends to hover around the critical value. The energy of the nanoflare, which combines the energies of the separate reconnections, does not have a wide variation. The energy range is typically less than 1.0 in the logarithm, i.e., factor of 10 difference between the largest and smallest for a given set of model parameters. Whether these assumptions are reasonable has yet to be determined. The detailed physical scenario of nanoflares is still unknown. We note that Ugarte-Urra \& Warren (2014) used a combination of observations and hydro modeling to infer a nanoflare repetition rate of 2 to 3 per hour. This corresponds to a delay of $\sim$1400 s, which agrees well with our assumed delay of 1000 s. From another observational/modeling comparison, Cargill (2014) suggests a delay of a few hundred to somewhat over 2000 s. Finally, Dahlburg et al. (2005) suggest approximately 2000 s based on the MHD simulations of the secondary instability. It is clear that much more work is left to be done.
16
7
1607.03883
We study a two-dimensional cellular automaton (CA) model for the evolution of coronal loop plasmas. The model is based on the idea that coronal loops are made of elementary magnetic strands that are tangled and stressed by the displacement of their footpoints by photospheric motions. The magnetic stress accumulated between neighbor strands is released in sudden reconnection events or nanoflares that heat the plasma. We combine the CA model with the Enthalpy Based Thermal Evolution of Loops model to compute the response of the plasma to the heating events. Using the known response of the X-Ray Telescope on board Hinode, we also obtain synthetic data. The model obeys easy-to-understand scaling laws relating the output (nanoflare energy, temperature, density, intensity) to the input parameters (field strength, strand length, critical misalignment angle). The nanoflares have a power-law distribution with a universal slope of -2.5, independent of the input parameters. The repetition frequency of nanoflares, expressed in terms of the plasma cooling time, increases with strand length. We discuss the implications of our results for the problem of heating and evolution of active region coronal plasmas.
false
[ "critical misalignment angle", "coronal loop plasmas", "strand length", "active region coronal plasmas", "field strength", "sudden reconnection events", "synthetic data", "photospheric motions", "elementary magnetic strands", "Loops model", "coronal loops", "neighbor strands", "board Hinode", "nanoflare energy", "the plasma cooling time", "heating", "nanoflares", "evolution", "laws", "the input parameters" ]
12.900204
15.360104
2
12437432
[ "Enea Romano, Antonio" ]
2016arXiv160708533E
[ "General background conditions for K-bounce and adiabaticity" ]
0
[ "-" ]
null
[ "astronomy" ]
5
[ "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "2001PhRvD..64l3522K", "2002PhRvD..65h6007K", "2004JHEP...05..074H", "2004PhRvD..69f3514E", "2005PhRvD..71b3515V", "2006JHEP...12..080C", "2007JCAP...11..010C", "2007JHEP...10..071C", "2007PhRvD..75h3513A", "2007PhRvD..76j3501L", "2007PhRvD..76l3503B", "2008JCAP...04..017L", "2008PhR...465..223L", "2010JCAP...10..026D", "2010PhRvL.104i1301K", "2011JCAP...11..021E", "2011PhRvD..84h3520X", "2012JCAP...08..020C", "2013PhRvD..87b3502D", "2014arXiv1406.3992P", "2016PhLB..755..464R", "2016PhLB..761..119R", "2016PhRvD..94d3511K", "2016arXiv160700996E", "2017JHEP...05..074A" ]
[ "10.48550/arXiv.1607.08533" ]
1607
1607.08533_arXiv.txt
The inflationary scenario provides a model able to explain many of the observational features of the observed Universe, such as the cosmic microwave background radiation isotropy or the spatial flatness. An alternative scenario to solve the horizon problem is the possibility that the Universe underwent a contraction phase before the big-bang, i.e. a bounce. Several attempts have been made to understand if the bounce could be a fully consistent cosmological model \cite{Gordon:2000hv,Khoury:2009my,Khoury:2001wf,Lehners:2008vx,Erickson:2003zm,Khoury:2001bz,Buchbinder:2007ad,ArkaniHamed:2003uy,Creminelli:2006xe,Lehners:2007ac,Creminelli:2007aq,Xue:2011nw,Cai:2012va,Cai:2007qw}, and could indeed be viable alternatives to inflation. In this paper we study the general conditions to realize the bounce with a K-essence type scalar field, without introducing any additional field. At the background level we impose the existence of two turning points where the derivative of the Hubble parameter $H$ changes sign and of a bounce point where the Hubble parameter vanishes. We first obtain some general conditions on the function $K(X)$ and the potential $V(\phi)$ and then give the examples of two models constructed according to these conditions. One is based on a quadratic $K(X)$, and the other on a $K(X)$ which is always avoiding divergences of the second time derivative of the scalar field, which may otherwise occur. % We integrate numerically the classical equations of motion for these models, verifying that they indeed have the expected evolution, and can produce a bounce. In the region where these models have a constant potential they became adiabatic on any scale, i.e. globally adiabatic, and because of this they may not conserve curvature perturbations on super-horizon scales. We then study the stability of the perturbations, focusingn the attention on the presence of ghosts and gradient instabilities. We confirm that gradient instabilities can arise around the bounce because the ghost free condition implies that the sign of the squared speed of sound is opposite to the sign of the time derivative of $H$. Finally we discuss how this kind of instabilities could be avoided by modifying the Lagrangian by introducing Galileion terms in order to prevent a negative squared speed of sound around the bounce.
We have derived the general conditions to obtain a bounce with a scalar field model with generalized kinetic term of the form $K(X)$ in general relativity. The requirement of the existence of a bounce point and of two turning points for $H(t)$ give some conditions which $K(X)$ and the potential $V(\phi)$ have to satisfy. We have given the examples of two models constructed according to these conditions. One is based on a quadratic $K(X)$, and the other on a $K(X)$ which is always avoiding divergences of the second time derivative of the scalar field, which may otherwise occur. An appropriate choice of the initial conditions can lead to a sequence of consecutive bounces. While at the perturbation level one class of models is free from ghosts, in general gradient instabilities are present around the bounce time, and they could be avoided only by modifying the Lagrangian. In the future it will be interesting to add Galileion type terms to the Lagrangian and study in details the cosmological perturbations during and after the bounce, and the effects of global adiabaticity in the region where the potential tends to a constant value. \appendix
16
7
1607.08533
We study the background conditions for a bounce uniquely driven by a single scalar field model with a generalized kinetic term $K(X)$, without any additional matter field. At the background level we impose the existence of two turning points where the derivative of the Hubble parameter $H$ changes sign and of a bounce point where the Hubble parameter vanishes. We find the conditions for $K(X)$ and the potential which ensure the above requirements. We then give the examples of two models constructed according to these conditions. One is based on a quadratic $K(X)$, and the other a $K(X)$ which is avoiding divergences of the second time derivative of the scalar field, which may otherwise occur. An appropriate choice of the initial conditions can lead to a sequence of consecutive bounces, or oscillations of $H$. In the region where these models have a constant potential they are adiabatic on any scale and because of this they may not conserve curvature perturbations on super-horizon scales. While at the perturbation level one class of models is free from ghosts and singularities of the classical equations of motion, in general gradient instabilities are present around the bounce time, because the sign of the squared speed of sound is opposite to the sign of the time derivative of $H$. We discuss how this kind of instabilities could be avoided by modifying the Lagrangian by introducing Galileion terms in order to prevent a negative squared speed of sound around the bounce.
false
[ "super-horizon scales", "consecutive bounces", "curvature perturbations", "models", "general gradient instabilities", "Galileion terms", "a single scalar field model", "sound", "instabilities", "the bounce time", "Hubble", "the second time derivative", "any additional matter field", "a negative squared speed", "a generalized kinetic term", "a bounce point", "the time derivative", "the scalar field", "order", "the squared speed" ]
10.544083
-0.269727
89
12473412
[ "Morford, J. C.", "Fenech, D. M.", "Prinja, R. K.", "Blomme, R.", "Yates, J. A." ]
2016MNRAS.463..763M
[ "e-MERLIN 21 cm constraints on the mass-loss rates of OB stars in Cyg OB2" ]
13
[ "Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK", "Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK", "Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK", "Royal Observatory of Belgium, Ringlaan 3, B-1180 Brussel, Belgium", "Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK" ]
[ "2017A&A...608A..69B", "2017AcASn..58...43A", "2017ApJ...845...39O", "2018A&A...617A.137F", "2018A&A...619A..54V", "2019A&A...627A..99N", "2019PASP..131a6001M", "2020A&A...637A..64M", "2020A&A...642A.168B", "2020ApJ...895...38H", "2022A&A...658A..61R", "2023AJ....166...23H", "2024ApJ...961..103D" ]
[ "astronomy" ]
3
[ "stars: early type", "stars: mass-loss", "galaxies: clusters: individual: (Cygnus OB2)", "radio continuum: stars", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1958ApJ...128...41S", "1975MNRAS.170...41W", "1984A&AS...57..327W", "1985ApJ...297..751D", "1988ApJ...335..914O", "1989ApJ...340..518B", "1989ApJS...71..267D", "1990ApJ...357..573D", "1990ApJ...361..607P", "1991AJ....101.1408M", "1991ApJ...377..629L", "1993ApJ...412..771L", "1994PASP..106.1025H", "1995A&A...299..523F", "1996A&A...311..264P", "1998A&A...332..251S", "1998Ap&SS.263..197T", "1998ApJ...496..407H", "1999A&A...350..181V", "2000A&A...354..193H", "2000A&A...360..539K", "2000A&A...362..295V", "2001A&A...366..508V", "2001A&A...366..623H", "2001A&A...369..574V", "2001ApJ...547..792D", "2002A&A...381.1015R", "2002A&A...382..921B", "2002A&A...382..999M", "2002A&A...389..874C", "2002A&A...396..949H", "2003A&A...408..715B", "2003ASPC..288..243H", "2003ApJ...597..957H", "2003ApJS..149..123S", "2004ESASP.552...33K", "2004PASP..116..909E", "2005A&A...435..239C", "2005A&A...436...17E", "2005A&A...437..657D", "2005A&A...440..261R", "2005A&A...441...23B", "2006A&A...446..279C", "2006A&A...454..625P", "2007A&A...464..211A", "2007ApJ...664.1102K", "2008A&A...478..823M", "2008A&A...481..777S", "2008A&A...487..575N", "2008A&ARv..16..209P", "2008cihw.conf...43N", "2009A&A...495..231L", "2009ARA&A..47..481A", "2010A&A...519A.111B", "2010A&A...521L..55P", "2010ApJ...713..871W", "2011A&A...526A..32M", "2011A&A...528A..64S", "2011A&A...535A..32N", "2011A&A...536A..31R", "2012A&A...537A..37M", "2012A&A...539A..79R", "2012A&A...541A.145C", "2012Ap.....55..371M", "2012ApJ...756...50K", "2013A&A...559A.130S", "2013A&C.....2...54P", "2013AstBu..68...87M", "2014A&A...561A..92C", "2014A&A...568A..59S", "2014AJ....147...40C", "2014MNRAS.438..639W", "2014MNRAS.439..908C", "2014ascl.soft07017A", "2015ApJ...811...85K", "2015ApJS..221....1R", "2015MNRAS.449..741W", "2016MNRAS.458..491M", "2016MNRAS.458.1999P" ]
[ "10.1093/mnras/stw1914", "10.48550/arXiv.1607.08828" ]
1607
1607.08828_arXiv.txt
The Cyg OB2 Radio Survey (COBRaS) is an e-MERLIN Legacy Project to carry out a deep imaging radio survey of the central region of the Cygnus OB2 association (www.merlin.ac.uk/legacy/projects/cobras.html). The principal component of this project (252 hours) will be to map the core of Cyg OB2 at 5\,GHz (C-band; 6 cm; 2 GHz full bandwidth), going to a depth of $\sim$ 3\,$\mu$Jy (1-$\sigmaup$). Prior to these observations which are due in late 2016, additional pointings (42 hrs) at 1.4\,GHz (L-band; 21cm; 512 MHz full bandwidth) have been secured during 2014. We report here on results from the supplementary L-band datasets, and specifically on the constraints they provide on the mass-loss rates of OB stars and the nature of their outer wind regions. Cygnus X is one of the richest star formation regions in the Galaxy. It hosts several OB associations, numerous young open clusters, tens of compact H\textsc{II} regions and star formation regions, a supernova remnant, and a superbubble blown by the collected stellar winds of the massive stars (e.g. \citealp{knodlseder_etal_2004}, \citealp{trapero_etal_1998}). At the core of Cygnus X is the Cyg OB2 association, which with a total cluster mass estimated to be $\sim$ 3$\times 10^4$\solmas, can be considered more as a massive cluster than an open OB association (\citealp{knodlseder_2000}; \citealp{wright_etal_2010}). The Cyg OB2 association is a uniquely important laboratory for studying the collective and individual properties of massive stars, and (possibly triggered) active star-formation. The stellar population of Cyg OB2 has been the focus of several studies across different wavebands (e.g. {\citealp{massey_thompson_1991}; \citealp{herrero_etal_2001}}; \citealp{comeron_etal_2002}; \citealp{setiagunawan_etal_2003}; \citealp{wright_etal_2014}; \citealp{rauw_etal_2015}). It has also been the target of radial velocity surveys (e.g. \citealp{kiminki_etal_2007}; \citealp{kobulnicky_etal_2012}). We lean here in particular on the recent census of \citet{wright_etal_2015} who list 169 OB stars, including 52 O-type and 8 normal early B supergiant. With an estimated cluster age of $\sim$ 2 Myr (\citealp{colombo_etal_2007}), Cyg OB2 is not only very rich in stellar density but also in its diversity. {The greater Cygnus X region includes} Be stars, many Young Stellar Objects (YSOs), two known Wolf-Rayet stars (WR 145, WR 146), two candidate Luminous Blue Variable (LBV) stars (G79.29+0.46, Cyg OB2 \#12), a red supergiant (IRC+40 427), a B[e] star (MWC 349), H\textsc{II} regions with groups of massive stars around them (DR 15, DR 18) and a gamma-ray source (TeV J2032+4130). Cyg OB2 is relatively close-by (at $\sim$ 1.4\,kpc), heavily obscured (as is the whole Cygnus X region), and located behind the Great Cygnus Rift. There is large and non-uniform visual extinction ranging from 4 to 10 mag (\citealp{knodlseder_2000}), thus making the association ideally studied at radio wavelengths. Second only to the initial stellar mass, the mass-loss rates of massive stars determine the final stellar mass, and thereby, the type of compact stellar remnant for all stars more massive than about 8 \solmas. The amount of mass shed during main- and post-main-sequence evolution determines whether a star becomes a black hole or a neutron star and specifies the type of supernova or gamma-ray burst that it may produce. Knowledge of stellar mass-loss remains one of the most uncertain parameters in massive star evolution because of the unknown amount of clumping in the stellar winds. {Results since the late 20th century have} strongly challenged the canonical model of stellar-wind mass-loss in massive stars by emphasising uncertainties in small-scale clumping and large-scale structure in the outflows. {In normal OB-type stars, small-scale clumps are optically thin in H$\alpha$ and most likely also in radio emission} leading to derived mass-loss rate ($\rm{\dot{M}}$) diagnostics that {have been found to disagree with one another} by a factor of {2 to 10 (e.g. \citealp{drew_1990};} \citealp{puls_etal_2006}; \citealp{prinja_massa_2010}; \citealp{muijres_etal_2011}). The observations indicate that the winds universally contain large structures and small-scale clumping that are only partially characterised observationally, with effects on mass-loss rates {that have yet to be fully understood in conjunction with one another (see \citealt{sundqvist_etal_2014}).} There is thus a pivotal requirement to constrain wind clumping as a function of radial distance/velocity from the surface of the star and make comparisons to theoretical predictions in order to derive reliable mass-loss rates. OB stars emit radio radiation through (thermal) free-free emission, due to electron-ion interactions in their ionised wind. The considerable advantage of using free-free radio fluxes for determining mass-loss for massive stars is that, unlike H$\alpha$ and UV, the emission arises at large radii in the stellar wind, where the terminal velocity will have been reached. The interpretation of the radio fluxes is more straightforward therefore and is not strongly dependent on details of the velocity law, ionisation conditions\footnote{Care must be taken when considering the ionisation state of He in the outer wind regions {for it has been known to alter inferred mass-loss rates \citep[see e.g.][]{lamers_leitherer_1993}.}}, inner velocity field, or the photospheric profile. Furthermore, the greater geometric region and density squared dependence of the free-free flux makes the radio observations very sensitive to clumping in the wind. The radio measurements can be directly compared to other density-squared diagnostics such as H$\alpha$, which in turn permits constraints on the relative amount of wind clumping as a function of velocity (see e.g. \citealp{blomme_etal_2002}; \citealp{puls_etal_2006}). We report here on first performance and science results from the COBRaS L-band (21 cm) Legacy data. We focus in this study on detection limits on thermal emission from suspected luminous single O and early B stars. The targets examined here are stars predicted to have the densest winds and highest mass-loss rates from inner-wind diagnostics such as H$\alpha$.
\label{discussion} We have used e-MERLIN observations of Cyg OB2 from the ongoing COBRaS project to demonstrate that the 21 cm flux densities of a sample of luminous, early O-type stars are below $\sim$ 70 $\mu$Jy. Under the assumption that the emission is entirely thermal in origin, and the stellar wind region is unclumped, we place upper limits of $\sim$ {4.4 $-$ 4.8} $\times$ 10$^{-6}$ M$_\odot$ yr $^{-1}$ on the mass-loss rates of O3 I stars; i.e. the hottest and most luminous stars in our sample. The mass-loss rates of early B supergiants (B0 to B1) are constrained to less than $\sim$ {1.8 $-$ 2.9} $\times$ 10$^{-6}$ M$_\odot$ yr $^{-1}$. Adopting spectroscopic masses, our upper limits are broadly consistent with mass-loss rates derived from the semi-empirical prescriptions of \citet{vink_etal_2000, vink_etal_2001}{, with the exception of the LBV candidate Cyg OB2 \#12.} For luminous O stars the \citealt{vink_etal_2001} values are in turn consistent with the refined predictions of \citealt{muijres_etal_2012}, who solve the wind dynamics numerically. The O3 to O5 stars in our sample have an effective photospheric radius of more than 150 R$_\star$ at 21 cm and our observations thus sample the most outer regions of the stellar winds. Assumptions that the wind is unclumped or very weakly clumped in this region is essentially untested observationally. Given that free-free emission depends on density squared, our fluxes either correspond to a smooth wind mass-loss rate, or a lower mass-loss rate $\times$ $\sqrt{f_{cl}}$, where $f_{cl}$ is the clumping factor. Comparing to the primarily recombination-formed line-synthesis analyses (i.e. H$\alpha$, He{\sc II}) of Cyg OB2 \#7, which samples the inner-most wind regions (below $\sim$ 3 R$_\star$), \citet{herrero_etal_2002}, \citet{mokiem_etal_2005}, \citet{repolust_etal_2005}, \citet{maryeva_etal_2013} all derive a 'smooth wind' mass-loss rate of $\sim$ 8.0 $-$ 10 $\times$ 10$^{-6}$ M$_\odot$ yr $^{-1}$. These consistently high mass-loss rates can only be reconciled with our 21 cm {\it upper limit} of {4.8} $\times$ 10$^{-6}$ \solmasyr if the inner wind H$\alpha$ region (close to the stellar surface) is substantially more clumped than the radio free-free formation region sampled in our study. This result is in agreement with the clumped wind models discussed by \citet{puls_etal_2006}, and with the notion that there is a radial stratification of the clumping factor in the stellar winds of OB stars. However, the derived clumping factor (and therefore mass-loss rate) is dependent on the assumption adopted for the degree of clumping in the radio formation region. Regarding the issue of structure in the outermost wind regions, the growth of the intrinsic line-deshadowing instability (LDI) has been numerically modelled by e.g. \citet{owocki_etal_1988}; \citet{feldmeier_1995}; \citet{dessart_owocki_2005}. The simulations show that the LDI leads to high-speed rarefactions that provide a basis for our interpretation of wind clumping. In their 1-d time-dependent hydrodynamical study of stochastic structure, Runacres {\&} Owocki (2002) model the evolution of clumped structure far from the stellar surface. Their models predict a rise in the clumping factor from the inner wind to $\sim$ 50 R$_\star$, and a subsequent decrease in the clumping factor to a residual value beyond $\sim$ 100 R$_\star$. Depending on the details, simulations predict that the stellar winds remain clumped deep into the radio formation region, with clumping factors between 2.5 to 6. As noted above, the single epoch radio continuum observations do not provide any direct information as to whether or not the OB stars winds are clumped beyond $\sim$ 100 R$_\star$. The substantial 6 cm (C-band) e-MERLIN COBRaS Legacy observations, scheduled from October 2016 onward, will provide flux densities down to a 3$\sigma$ limit of $\sim$ 10 $\mu$Jy. These data will ultimately lead to the tightest constraints on the outer wind mass-loss rates of OB stars in Cyg OB2 for a wide range of effective temperature, luminosity and wind density.
16
7
1607.08828
We present e-MERLIN 21 cm (L-band) observations of single luminous OB stars in the Cygnus OB2 association, from the Cyg OB2 Radio Survey Legacy programme. The radio observations potentially offer the most straightforward, least model-dependent, determinations of mass-loss rates, and can be used to help resolve current discrepancies in mass-loss rates via clumped and structured hot star winds. We report here that the 21 cm flux densities of O3 to O6 supergiant and giant stars are less than ∼70 μJy. These fluxes may be translated to `smooth' wind mass-loss upper limits of ∼4.4-4.8 × 10<SUP>-6</SUP> M<SUB>⊙</SUB> yr <SUP>-1</SUP> for O3 supergiants and ≲2.9 × 10<SUP>-6</SUP> M<SUB>⊙</SUB> yr <SUP>-1</SUP> for B0 to B1 supergiants. The first ever resolved 21 cm detections of the hypergiant (and luminous blue variable candidate) Cyg OB2 #12 are discussed; for multiple observations separated by 14 d, we detect an ∼69 per cent increase in its flux density. Our constraints on the upper limits for the mass-loss rates of evolved OB stars in Cyg OB2 support the model that the inner wind region close to the stellar surface (where Hα forms) is more clumped than the very extended geometric region sampled by our radio observations.
false
[ "structured hot star winds", "single luminous OB stars", "Cyg OB2", "evolved OB stars", "giant stars", "Cygnus OB2", "multiple observations", "luminous blue variable candidate", "∼70 μJy", "O6 supergiant", "O3 supergiants", "Hα forms", "the Cyg OB2 Radio Survey Legacy programme", "mass-loss rates", "B1 supergiants", "the Cyg OB2 Radio Survey Legacy", "Cyg", "current discrepancies", "the Cygnus OB2 association", "`smooth wind mass-loss upper limits" ]
9.05359
10.670706
153
12518439
[ "Ghelfi, A.", "Maurin, D.", "Cheminet, A.", "Derome, L.", "Hubert, G.", "Melot, F." ]
2017AdSpR..60..833G
[ "Neutron monitors and muon detectors for solar modulation studies: 2. ϕ time series" ]
50
[ "LPSC, Université Grenoble-Alpes, CNRS/IN2P3, 53 avenue des Martyrs, 38026 Grenoble, France", "LPSC, Université Grenoble-Alpes, CNRS/IN2P3, 53 avenue des Martyrs, 38026 Grenoble, France", "ONERA (French Aerospace Lab), 2 avenue Edouard Belin, 31055 Toulouse Cedex 4, France", "LPSC, Université Grenoble-Alpes, CNRS/IN2P3, 53 avenue des Martyrs, 38026 Grenoble, France", "ONERA (French Aerospace Lab), 2 avenue Edouard Belin, 31055 Toulouse Cedex 4, France", "LPSC, Université Grenoble-Alpes, CNRS/IN2P3, 53 avenue des Martyrs, 38026 Grenoble, France" ]
[ "2016PhRvD..94l3007F", "2016PhRvD..94l3019K", "2017ApJ...844L..26T", "2017ICRC...35..271T", "2017JGRA..122.3875U", "2017JGRA..122.9790A", "2017JGRA..12210964G", "2017JInst..12P2021L", "2017PhRvD..96d3007D", "2017PhRvD..96j3005T", "2017PhRvL.119x1101G", "2017ehep.confE.620T", "2018ApJ...858...43M", "2018ApJ...863..119Z", "2018JGRA..123.1731S", "2018JGRA..123.7181N", "2018PhRvD..97j3011K", "2019ApJ...883...73C", "2019ApJ...887..132S", "2019ITNS...66..262H", "2019JGRA..124..661H", "2019PhRvD..99b3016P", "2020A&A...639A..74W", "2020A&A...639A.131W", "2020ApJ...903...21M", "2020ITAES..56.3674H", "2020JGRA..12527304N", "2020Univ....6..102M", "2021APh...12402495Z", "2021AdSpR..68.2974M", "2021ApJ...921..109S", "2021JCAP...03..099D", "2021JCAP...07..010D", "2021JCAP...11..018D", "2021PhRvD.104h3005G", "2021PhRvD.104l3001Z", "2021SpWea..1902665H", "2022A&A...667A..25M", "2022ApJ...932...37N", "2022PhRvD.106f3006W", "2022ScPP...12..163C", "2022arXiv220306479V", "2023EPJC...83..971M", "2023JCAP...10..011D", "2023MNRAS.518.4832F", "2023RAA....23h5019X", "2023Univ....9..502W", "2024APh...15902949H", "2024JCAP...05..104D", "2024PhRvD.109b3023L" ]
[ "astronomy" ]
7
[ "Solar modulation", "Cosmic rays", "Neutron monitor", "Muon detector", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1956PhRv..104.1723M", "1957PMag....2..157F", "1957PhRv..107.1386M", "1959PhRv..115..194M", "1959PhRv..116..462M", "1960JGR....65..767M", "1961PhRv..121.1206A", "1963PhRv..129.2275M", "1964CaJPh..42.2443H", "1964JGR....69.3097W", "1964JGR....69.3293F", "1964JGR....69.3939S", "1964PhRv..133..818F", "1964PhRvL..13..106O", "1965JGR....70.2111F", "1965JGR....70.3515F", "1965JGR....70.5753F", "1966JGR....71.1771B", "1966P&SS...14..503C", "1966PhRvL..16..813F", "1967ApJ...149L.115G", "1967JGR....72.2765D", "1967P&SS...15..715H", "1967PhRv..163.1327B", "1968ApJ...154.1011G", "1968JGR....73.4231O", "1968JGR....73.4261F", "1969ApJ...155..609C", "1971ApJ...166..221H", "1971JGR....76.1605L", "1971JGR....76.7445R", "1972PhRvL..28..985R", "1973ApJ...180..987S", "1973ICRC....2..732G", "1974JGR....79.4127R", "1975ApJ...202..265G", "1976ApJ...206..616M", "1978ApJ...221.1110L", "1979ICRC....1..330B", "1983ApJ...275..391W", "1985ApJ...294..455B", "1986ApJ...303..816K", "1987ICRC....1..325W", "1989NCimC..12..173N", "1990ICRC....7...96N", "1991ApJ...378..763S", "1991ApJ...380..230W", "1991NCimC..14..213C", "1993ApJ...413..268B", "1995ICRC....2..630W", "1997NIMPA.389...81B", "1999ApJ...518..457B", "1999CzJPh..49.1743U", "1999ICRC....7..317C", "1999PhRvD..60e2002B", "2000ApJ...533..281M", "2000JGR...10527447B", "2000PhLB..490...27A", "2000PhLB..494..193A", "2000SSRv...93...11S", "2000SSRv...93..335C", "2000SoPh..195..175C", "2001AdSpR..26.1831S", "2002ApJ...564..244W", "2002SoPh..207..389U", "2003APh....19..583B", "2003ApJ...591.1220L", "2003JGRA..108.1355W", "2004ASSL..303.....D", "2004JGRA..109.1101C", "2005JGRA..11012108U", "2007APh....28..154S", "2007ARNPS..57..285S", "2008AdSpR..42..510D", "2008ICRC....1..289F", "2008ICRC....1..733S", "2008ICRC....1..737S", "2009ASSL..358.....D", "2009AdSpR..44.1107S", "2009JGRA..114.8104M", "2011APh....34..401D", "2011AdSpR..48...19M", "2011ICRC...11..467A", "2011JGRA..116.2104U", "2011JInst...6P1003P", "2011Sci...332...69A", "2012AdSpR..49.1563D", "2012CRPhy..13..740L", "2012ITNS...59.1722C", "2012JGRA..11712103C", "2012JInst...7C4007C", "2013A&ARv..21...70B", "2013AdSpR..51..219A", "2013ApJ...765...91A", "2013ITNS...60.2411C", "2013ITNS...60.2418H", "2013JGRA..118.2783M", "2013JGRA..118.7488C", "2013LRSP...10....1U", "2014A&A...569A..32M", "2015AdSpR..55..363M", "2015AdSpR..55.2940U", "2015JGRA..120.7172G", "2015NIMPA.798..172P", "2015PhRvL.114q1103A", "2015PhRvL.115u1101A", "2015nfmw.book.....L", "2016A&A...591A..94G", "2016Ap&SS.361...48B", "2016ApJ...818...68A", "2016ApJ...822...65A", "2016SoPh..291..965K" ]
[ "10.1016/j.asr.2016.06.027", "10.48550/arXiv.1607.01976" ]
1607
1607.01976.txt
Measurements of top-of-atmosphere (TOA) cosmic-ray (CR) fluxes show a clear modulation related to solar activity \citepads{2013LRSP...10....1U}. The imprint of the 11-year solar cycle is present in secondary particles created in the Earth atmosphere \citepads{1974crvs.book.....D,2004ASSL..303.....D,2009crme.book.....D}, as seen in neutron monitor data \citepads{2000SSRv...93...11S}. Despite being an integral measurement (top-of-atmosphere fluxes folded by the atmosphere and instrument response), ground-based detectors have been providing monitoring of solar activity since the 50's, on a much finer timescale than balloon-borne and space experiments can achieve, even today. In this work, we wish to provide a consistent description of modulation levels for cosmic-ray data. This is important in the context of galactic CR physics as clues on CR sources \citepads{2013A&ARv..21...70B} and constraints set on CR transport parameters \citepads{2007ARNPS..57..285S} are based on modulated CR data. Similarly, dark matter indirect detection \citepads{2012CRPhy..13..740L} involves low energy modulated antiproton and antideuteron fluxes. Unfortunately, the set of modulation levels provided for space or balloon-borne CR data (from the original publications) is not homogeneous and very patchy \citepads{2014A&A...569A..32M}: each value, when existing, is based on different assumptions regarding the IS spectrum (fitted to the experiment data, or resulting from different CR propagation models) and the modulation model (from force-field to sign-charge dependent drift models). This situation is inadequate and unsatisfactory. Providing homogeneous modulation levels for past and present CR experiments or providing $\phi(t)$ time series are complementary tasks. In the context of the force-field approximation \citepads{1967ApJ...149L.115G,1968ApJ...154.1011G}, homogeneous monthly time series have been derived from NM data \citepads{1999CzJPh..49.1743U,2002SoPh..207..389U,2005JGRA..11012108U,2011JGRA..116.2104U} since July 1936. Note however, that many experiments operate on a shorter timescale, during which solar activity can significantly depart from the monthly average. This is especially true during solar maximum periods. Moreover, in the last years, the PAMELA\footnote{\url{http://pamela.roma2.infn.it}.} \citepads{2011Sci...332...69A,2013AdSpR..51..219A,2013ApJ...765...91A} and AMS\footnote{\url{http://www.ams02.org}.} \citepads{2015PhRvL.114q1103A,2015PhRvL.115u1101A} experiments provided high precision proton and helium fluxes. The latter can be used to improve the IS flux description \citepads{2016Ap&SS.361...48B,2015arXiv151108790C,2016A&A...591A..94G}, and in a second step the accuracy of $\phi$ time series. Our approach is based on \citetads{2011JGRA..116.2104U}, with several differences. We build on our recent analysis of the uncertainties on $\phi$ reconstruction from ground-based detectors data \citepads[][hereafter Paper~I]{2015AdSpR..55..363M}. We also take advantage of our recent re-estimate of the interstellar (IS) proton and helium fluxes \citepads{2016A&A...591A..94G}. The robustness and consistency of $\phi$ time series from NM data (retrieved from the Neutron Monitor Data Base, NMDB\footnote{\url{http://www.nmdb.eu}.}) are validated against GCR data $\phi$ values and compared to other ground-based detector data. The paper is organised as follows. In Sect.~\ref{sec:detectors}, we recall how IS spectra are modulated and folded by the yield function of ground-based detectors, whose data are used to reconstruct $\phi$ time-series. In Sect.~\ref{sec:sZsup2}, we discuss the enhancement factor to account for heavy CR contributions to count rates. In Sect.~\ref{sec:nm_norm}, the procedure to calculate the correction factor of a NM station is detailed. In Sect.~\ref{sec:time-series}, we calculate and compare $\phi$ time-series (and their uncertainties) as obtained from NM, GCR, Auger scaler, or neutron spectrometer data. We conclude in Sect.~\ref{sec:concl}. Along with the paper, we provide an online application to calculate, at any time in the past, $\phi$ values (for any time period) based on the methodology presented in this paper. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\label{sec:concl} We have revisited and extended the analysis of \citetads{2011JGRA..116.2104U} to refine the calculation of $\phi$ time series (and uncertainties) from three type of ground-based detectors (NM, BSS and $\mu$-like detectors): \begin{itemize} \item Our analysis benefits from the improvements made on the determination of the H and He IS fluxes and their uncertainties \citepads{2016A&A...591A..94G}. The associated error $\Delta\phi_{J^{\rm IS}}$ is estimated to be no more than $\pm25$~MV for NMs and BSS, and $\pm18$~MV for $\mu$-like detectors, to be compared to the at-the-time conservative 200~MV of Paper~I. \item A common assumption to calculate count rates is to fold the H and He TOA fluxes by the yield function, accounting for $Z>2$ CRs as an enhancement factor $s_{Z>2}$ for He. Following the minute approach of \citetads{2016A&A...591A..94G} to extract $Z>2$ IS fluxes, we find $s_{Z>2}=0445\pm 0.03$. This uncertainty comes from the assumption of a rigidity-independent enhancement, which at present provides a negligible contribution to the total error budget of $\phi^{\rm NM,\,BSS,\,\mu}(t)$. Improvements in the calculation of $\phi^{\rm NM}$ will however require to use an energy-dependent scaling to keep this uncertainty subdominant. \item The uncertainty from the yield function is estimated from the scatter obtained when using different parametrisations, leading to $(\sigma_\phi/\phi)_y\approx 6\%$ for NMs and BSS (negligible for $\mu$). Actually, most of the yield function uncertainty is absorbed in the correction factor $k^{\rm corr}$ that must be estimated for each detector (to account for environment effects). The spread in count rates from the scatter of the yield function, taken at face value, leads to a 25\% scatter on $\phi$ (as estimated in Paper~I), but gives the above 6\% spread on $\phi^{\rm NM,\,BSS}$ when all the calculation steps are carried out. \item A key step is the calculation of $k^{\rm corr}$, which is tackled by the use of carefully selected TOA CR data on which to normalise the detector. This allows us to reduce the uncertainty to $\Delta k^{\rm corr}/k^{\rm corr}=\pm2.2\%$ ($\pm3.6\%$ for $\mu$), leading to a $\pm50$~MV uncertainty for NM and BSS, and $\pm310$~MV for $\mu$. This is the dominant source of uncertainty, by far for $\mu$-like detectors, but it could be decreased as more TOA measurements become available. \end{itemize} This analysis shows that with an improved IS flux description, $\phi$ values extracted from NM data are now in agreement with those extracted from TOA data|a similar conclusion was recently reached by \citetads{2015AdSpR..55.2940U} analysing a Forbush decrease. The two sets of data are complementary: NM count rate data have a good time sampling and a stable setup over time, which are useful properties to reconstruct robust $\phi$ time series; CR TOA data, though scarcer and sometimes suffering from systematics, provide differential fluxes with increasing precision (as new instruments and techniques are used) that help to properly calibrate IS fluxes and $\phi$ time series. Further improvements of $\phi^{\rm NM}(t)$ time-series would require to take properly into account the Earth geomagnetic field and its long term evolution (for a better description of the transmission function). Improving on the yield function parametrisation and taking advantage of ongoing TOA measurements to further reduce the IS flux uncertainties is also desired. For the other detector types, BSS data show promising potential for intercalibration with NM data, in particular their complementarity to study snow falls effects on the neutron spectrum and count rates. Modulation time-series from $\mu$ detectors suffer from larger uncertainties than those from NM (due to the precision of TOA data used to calibrate their efficiency), but they also are complementary as they are not sensitive to the same uncertainties. In this respect, future Auger data from their histogram mode ($\mu$-like counter) are awaited to further investigate this complementarity. To conclude, we refer the interested reader to \url{http://lpsc.in2p3.fr/crdb}, where we provide an online tool to extract $\phi^{\rm NM}(t)$ time series, and/or the average $\langle\phi^{\rm NM}\rangle_{\Delta t}$ on a given time interval, and/or modulated fluxes, based on this analysis. The new $\phi^{\rm NM}(t)$ values are also used to provide homogeneous sets of values for all CR data in CRDB.
16
7
1607.01976
The level of solar modulation at different times (related to the solar activity) is a central question of solar and galactic cosmic-ray physics. In the first paper of this series, we have established a correspondence between the uncertainties on ground-based detectors count rates and the parameter ϕ (modulation level in the force-field approximation) reconstructed from these count rates. In this second paper, we detail a procedure to obtain a reference ϕ time series from neutron monitor data. We show that we can have an unbiased and accurate ϕ reconstruction (Δϕ / ϕ ≃ 10 %). We also discuss the potential of Bonner spheres spectrometers and muon detectors to provide ϕ time series. Two by-products of this calculation are updated ϕ values for the cosmic-ray database and a web interface to retrieve and plot ϕ from the 50's to today (http://lpsc.in2p3.fr/crdb).
false
[ "ϕ time series", "solar modulation", "neutron monitor data", "different times", "rates", "a reference ϕ time series", "detectors", "the parameter ϕ (modulation level", "today", "solar and galactic cosmic-ray physics", "an unbiased and accurate ϕ reconstruction", "these count rates", "ground-based detectors", "(Δϕ / ϕ", "the solar activity", "the cosmic-ray database", "a web interface", "spectrometers", "a central question", "the force-field approximation" ]
6.767267
0.038907
13
12409071
[ "Nakajima, Tadashi", "Sorahana, Satoko" ]
2016ApJ...830..159N
[ "Carbon-to-oxygen Ratios in M Dwarfs and Solar-type Stars" ]
9
[ "Astrobiology Center, 2-21-1, Osawa, Mitaka, Tokyo, 181-8588, Japan ;;", "Department of Astronomy, Graduate School of Science, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo, 113-0033, Japan" ]
[ "2017PASJ...69....1T", "2018MNRAS.476.3432P", "2019A&A...630A.104A", "2019A&A...630A.152H", "2019AsBio..19..867H", "2022A&A...661A.109D", "2023ApJ...944...41I", "2023MNRAS.521.5761G", "2024ApJ...961..210P" ]
[ "astronomy" ]
6
[ "stars: abundances", "stars: low-mass", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1987AJ.....93..864J", "1989GeCoA..53..197A", "1995AJ....110.1838R", "1996AJ....112.2799H", "2000Icar..145..637G", "2002A&A...390..235N", "2003Msngr.114...20M", "2005ApJ...620..481F", "2005ApJS..159..141V", "2005PASJ...57...65T", "2005astro.ph..4214K", "2009ARA&A..47..481A", "2009ApJ...691L.119S", "2010A&A...514A..92C", "2010ApJ...715.1050B", "2010ApJ...725.2349D", "2010JQSRT.111.2139R", "2011AJ....141...97W", "2011ApJ...735...41P", "2012ApJ...747L..27F", "2013A&A...552A..73N", "2013A&A...553A...6H", "2013AJ....145..102L", "2014A&A...568A..25N", "2014MNRAS.443.2561G", "2014PASJ...66...98T", "2015ApJ...804...40G", "2015PASJ...67...26T", "2016MNRAS.455.3824G", "2016PASJ...68...13T" ]
[ "10.3847/0004-637X/830/2/159", "10.48550/arXiv.1607.07528" ]
1607
1607.07528_arXiv.txt
The study of elemental abundances in the stellar atmosphere is an important part of modern astronomy. Carbon and oxygen are among the most important elements and the carbon-to-oxygen ratio is expected to be the essential parameter of the atmosphere of the star hosting planets, which controls the nature of the planets. In the early 2000's, possible existence of carbon dominated planets with carbon dominated silicates was discussed \citep{gai00,kuc05}. In the early 2010's, some determinations of the carbon-to-oxygen ratios in solar-type stars found that $\sim$25\% of nearby stars have C/O $>$ 0.8 and $\sim$8\% have C/O $>$ 1.0 \citep{bon10,del10,pet11}. \citet{bon10} suggested that the mineralogy of planets formed in the environmental gas with C/O $>0.8$ will make carbon dominated rocky planets, rather than oxygen dominated planets. These large fractions of high C/O ratio stars were criticized by \citet{for12} who pointed out some issues. { First the frequency of high C/O is inconsistent with relatively small fraction of dwarf carbon stars ($<10^{-3}$) in large samples of low-mass stars}. Second, these high C/O ratios are overestimated. { The possible reasons for this overestimation are a high C/O ratio for the Sun used for differential abundance analysis { and} the treatment of a Ni blend that affects the O abundance. } A careful analysis of carbon and oxygen abundances of FG main sequence stars in the solar neighborhood was presented by \citet{nis14} by taking into account the non-LTE effects in the model atmosphere and a critical analysis of the Ni blend to the forbidden [OI] line at 6300\AA. Among 66 nearby disk stars for which $\lambda$7774 OI triplet was analyzed with the non-LTE analysis, no star with C/O$>0.8$ was found. { \citet{nis14} conclude that C/O does not exceed 0.8, and this result seems to exclude formation of carbon dominated planets.} Star formation history of M dwarfs is expected not to be so different from that of solar-type stars and they can also be used as probes of high C/O in the solar neighborhood. The spectra of M dwarfs include absorption bands of O-containing molecules such as TiO, VO, and CaOH, which are sensitive to the available O abundance, and hence C/O. In a two dimensional space span by TiO and CaH indices, \citet{gai15} and \citet{giz16} plot a sequence of M dwarfs from which a small number of stars deviate in the direction of higher C/O. \citet{gai15} estimates that high C/O$\sim 1$ stars constitute less than $6 \times 10^{-4}$ of M dwarfs with 99\% confidence, while \citet{giz16} show that a high C/O ratio (0.9) is less than 1\%. Recently carbon and oxygen abundances of 46 M dwarfs have been obtained respectively from CO and H$_2$O molecular lines { derived from $K$ band spectra obtained at the Subaru telescope with a resolution of 20,000 } \citep{tsu14,tsu15,tsu16}. In this paper, we present the C/O ratios of these M dwarfs and examine the kinematics of the sample in terms of stellar population in \S2 and compare the distribution of C/O of these M dwarfs with those of solar-type stars obtained by others using the Kolmogorov-Smirnov test in \S3. Implications of the comparison results are discussed in \S4.
Recently carbon and oxygen abundances of 46 M dwarfs have been obtained respectively from CO and H$_2$O molecular lines \citep{tsu14,tsu15,tsu16}. This is not a differential abundance analysis with respect to the Sun. We present C/O ratios and kinematics of these M dwarfs. The distribution of C/O ratios in M dwarfs is compared with those in solar type stars obtained by \citet{nis14}, \citet{del10}, and \citet{pet11} using the K-S test. The distribution C/O ratios in M dwarfs and that by \citet{nis14} are consistent with being drawn from a same distribution, while the distributions by \citet{del10} and \citet{pet11} are not drawn from the same distribution as that in M dwarfs. High solar C/O ratios adopted by \citet{del10} and \citet{pet11} partly explain these results. Since carbon and oxygen abundances of solar type stars have been obtained by differential analyses with respect to the Sun, pairs of C/O ratio distributions with respect to the solar ratio, are compared among those by \citet{tak05}, \citet{nis14}, \citet{del10} and \citet{pet11}. { The distributions by \citet{tak05} and \citet{nis14} are consistent and those by \citet{del10} and \citet{pet11} are barely consistent, while other pairs are mutually inconsistent.} Larger scatters in C/O distributions by \citet{del10} and \citet{pet11} are explained by the use of [OI] 6300 forbidden line as the abundance indicator of oxygen.
16
7
1607.07528
It has been suggested that high C/O ratios (&gt;0.8) in circumstellar disks lead to the formation of carbon-dominated planets. Based on the expectation that elemental abundances in the stellar photospheres give the initial abundances in the circumstellar disks, the frequency distributions of C/O ratios of solar-type stars have been obtained by several groups. The results of these investigations are mixed. Some find C/O &gt; 0.8 in more than 20% of stars, and C/O &gt; 1.0 in more than 6%. Others find C/O &gt; 0.8 in none of the sample stars. These works on solar-type stars are all differential abundance analyses with respect to the Sun and depend on the adopted C/O ratio in the Sun. Recently, a method of molecular line spectroscopy of M dwarfs, in which carbon and oxygen abundances are derived respectively from CO and H<SUB>2</SUB>O lines in the K band, has been developed. The resolution of the K-band spectrum is 20,000. Carbon and oxygen abundances of 46 M dwarfs have been obtained by this nondifferential abundance analysis. Carbon-to-oxygen ratios in M dwarfs derived by this method are more robust than those in solar-type stars derived from neutral carbon and oxygen lines in the visible spectra because of the difficulty in the treatment of oxygen lines. We have compared the frequency distribution of C/O distributions in M dwarfs with those of solar-type stars and have found that the low frequency of high-C/O ratios is preferred.
false
[ "oxygen lines", "differential abundance analyses", "stars", "elemental abundances", "M dwarfs", "circumstellar disks", "molecular line spectroscopy", "oxygen", "high C/O ratios", "M", "Carbon and oxygen abundances", "carbon and oxygen abundances", "neutral carbon and oxygen lines", "solar-type stars", "several groups", "C/O ratios", "Sun", "gt", "this nondifferential abundance analysis", "the adopted C/O ratio" ]
8.095257
13.515248
-1
11453947
[ "Eichhorn, Helge", "Anderl, Reiner" ]
2016arXiv160700849E
[ "Plyades: A Python Library for Space Mission Design" ]
0
[ "-", "-" ]
null
[ "astronomy" ]
8
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
null
[ "10.48550/arXiv.1607.00849" ]
1607
1607.00849_arXiv.txt
% } Designing a space mission trajectory is a computation-heavy task. Software tools that conduct the necessary numerical simulations and optimizations are therefore indispensable and high numerical performance is required. No science mission or spacecraft are exactly the same and during the early mission design phases the technical capabilities and constraints change frequently. Therefore high development speed and programmer productivity are required as well. Due to its supreme numerical performance Fortran has been the top programming language in many organizations within the astrodynamics community. At the European Space Operations Center (ESOC) of the European Space Agency (ESA) a large body of sometimes decades old Fortran77 code remains in use for mission analysis tasks. While this legacy code mostly fulfills the performance requirements\DUfootnotemark{id1}{id3}{1} usability and programmer productivity suffer. One reason for this is that Fortran is a compiled language which hinders exploratory analyses and rapid prototyping. The low level of abstraction supported by Fortran77 and programming conventions, like the 6-character variable name limit and fixed-form source code, that are in conflict with today's best practices are a more serious problem, though. The more recent Fortran standards remedy a lot of these shortcomings, e.g. free-form source in Fortran90 or object-oriented programming features in Fortran2003, but also introduce new complexity, e.g. requiring sophisticated build systems for dependency resolution. Compiler vendors have also been very slow to implement new standards. For example this year the Intel Fortran compiler achieved full support of the Fortran2003 standard, which was released in 2005 \cite{IFC15}.% \DUfootnotetext{id3}{id1}{1}{ Many routines were not written with thread-safety and re-entrancy in mind and can therefore not be used in parallel codes.} Due to these reasons Fortran-based tools and libraries have been generally used together with programming environments with better usability such as MATLAB. A common approach for developing mission design software at ESOC is prototyping and implementing downstream processes such as visualization in MATLAB and then later porting performance-intensive parts or the whole system to Fortran77. The results are added complexity through the use of the MEX-interface for integrating Fortran and MATLAB, duplicated effort for porting, and still a low-level of abstraction because the system design is constrained by Fortran's limitations. Because of the aforementioned problems some organizations explore possibilities to replace Fortran for future development. The French space agency CNES (Centre National D'Études Spatiales) for instance uses the Java-based Orekit library \cite{Ore15} for its flight dynamics systems. In this paper we show why Python and the scientific Python ecosystem are a viable choice for the next generation of space mission design software and present the Plyades library. Plyades is a proof-of-concept implementation of an object-oriented astrodynamics library in pure Python. It makes use of many prominent scientific Python libraries such as Numpy, Scipy, Matplotlib, Bokeh, and Astropy. In the following we discuss the design of the Plyades data model and conclude the paper with an exemplary analysis.
% } In this paper we have discussed the current tools and programming environments for space mission design. These suffer from high complexity, low levels of abstraction, low flexibility, and out-dated programming practices. We have then shown why the maturity and breadth of the scientific Python ecosystem as well as the usability of the Python programming language make Python a viable alternative for next generation astrodynamics tools. With the design and implementation of the proof-of-concept library Plyades we demonstrated that it is possible to create powerful yet simple to use astrodynamics tools in pure Python by using scientific Python libraries and following modern best practices. The Plyades work has lead to the foundation of the Python Astrodynamics project, an inter-european collaboration, whose goal is the development of a production-grade Python-based astrodynamics library.
16
7
1607.00849
Plyades: A Python Library for Space Mission Design Designing a space mission is a computation-heavy task. Software tools that conduct the necessary numerical simulations and optimizations are therefore indispensable. The usability of existing software, written in Fortran and MATLAB, suffers because of high complexity, low levels of abstraction and out-dated programming practices. We propose Python as a viable alternative for astrodynamics tools and demonstrate the proof-of-concept library Plyades which combines powerful features with Pythonic ease of use.
false
[ "Pythonic ease", "use", "Space Mission Design Designing", "low levels", "powerful features", "high complexity", "astrodynamics tools", "Software tools", "Pythonic", "Plyades", "optimizations", "concept", "existing software", "MATLAB", "a computation-heavy task", "abstraction and out-dated programming practices", "Fortran", "the necessary numerical simulations", "Python", "a space mission" ]
10.327785
4.368825
170
12460181
[ "Lovelace, Geoffrey", "Lousto, Carlos O.", "Healy, James", "Scheel, Mark A.", "Garcia, Alyssa", "O'Shaughnessy, Richard", "Boyle, Michael", "Campanelli, Manuela", "Hemberger, Daniel A.", "Kidder, Lawrence E.", "Pfeiffer, Harald P.", "Szilágyi, Béla", "Teukolsky, Saul A.", "Zlochower, Yosef" ]
2016CQGra..33x4002L
[ "Modeling the source of GW150914 with targeted numerical-relativity simulations" ]
76
[ "Gravitational Wave Physics and Astronomy Center, California State University Fullerton, Fullerton, CA 92834, USA ;;", "Center for Computational Relativity and Gravitation, School of Mathematical Sciences, Rochester Institute of Technology, 85 Lomb Memorial Drive, Rochester, NY, 14623, USA", "Center for Computational Relativity and Gravitation, School of Mathematical Sciences, Rochester Institute of Technology, 85 Lomb Memorial Drive, Rochester, NY, 14623, USA", "Theoretical Astrophysics 350-17, California Institute of Technology, Pasadena, CA 91125, USA", "Gravitational Wave Physics and Astronomy Center, California State University Fullerton, Fullerton, CA 92834, USA", "Center for Computational Relativity and Gravitation, School of Mathematical Sciences, Rochester Institute of Technology, 85 Lomb Memorial Drive, Rochester, NY, 14623, USA", "Center for Astrophysics and Planetary Science, Cornell University, Ithaca, NY 14853, USA", "Center for Computational Relativity and Gravitation, School of Mathematical Sciences, Rochester Institute of Technology, 85 Lomb Memorial Drive, Rochester, NY, 14623, USA", "Theoretical Astrophysics 350-17, California Institute of Technology, Pasadena, CA 91125, USA", "Center for Astrophysics and Planetary Science, Cornell University, Ithaca, NY 14853, USA", "Canadian Institute for Theoretical Astrophysics, 60 St. George Street, University of Toronto, Toronto, ON M5S 3H8, Canada", "Theoretical Astrophysics 350-17, California Institute of Technology, Pasadena, CA 91125, USA; Caltech JPL, Pasadena, CA 91109, USA", "Center for Astrophysics and Planetary Science, Cornell University, Ithaca, NY 14853, USA", "Center for Computational Relativity and Gravitation, School of Mathematical Sciences, Rochester Institute of Technology, 85 Lomb Memorial Drive, Rochester, NY, 14623, USA" ]
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[ "astronomy", "physics" ]
3
[ "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.1088/0264-9381/33/24/244002", "10.48550/arXiv.1607.05377" ]
1607
1607.05377_arXiv.txt
\label{sec:intro} On September 14, 2015, the Advanced Laser Interferometer Gravitational-Wave Observatory (Advanced LIGO) directly measured gravitational waves for the first time~\cite{Abbott:2016blz}, giving birth to a new era of astronomy. The waves were emitted by a pair of black holes with masses $36_{-4}^{+5}$\,$M_{\odot}$\ and $29_{-4}^{+4}$\,$M_{\odot}$~\cite{Abbott:2016blz,TheLIGOScientific:2016wfe} that orbited each other, collided, and merged about 1.3 billion years ago. The gravitational wave signal from this event, named GW150914, is consistent with general relativity~\cite{TheLIGOScientific:2016src}. LIGO has already observed a second gravitational-wave signal from merging black holes (called GW151226)~\cite{Abbott:2016nmj} and a third possible signal (called LVT151012)~\cite{TheLIGOScientific:2016pea}; many more such observations are expected soon. Extrapolating from the observations of GW150914 and GW151226 and including LVT151012, Advanced LIGO is expected to observe between roughly 5 to tens of gravitational-wave signals from merging black holes during its next six-month observing run (O2) starting in 2016~\cite{TheLIGOScientific:2016pea}. The GW150914 observation included 8 gravitational-wave cycles, covering the late inspiral, merger, and ringdown phases of the binary (cf. Fig.~2 of Ref.~\cite{Abbott:2016blz}); this late phase of a binary-black-hole (BBH) merger can be described accurately only by directly solving the full equations of general relativity. To extract and validate robust conclusions about the astrophysical and fundamental significance of these events \cite{Abbott:2016blz,TheLIGOScientific:2016wfe,GW150914-ASTRO,TheLIGOScientific:2016src}, correct and complete solutions to Einstein's equations will be critical, and can be obtained only through direct numerical simulation. The first attempts to solve the general relativity field equations numerically date to the 1960s, when Hahn and Lindquist~\cite{HahnLindquist1964} performed the first studies. Smarr followed these early efforts with some success in the 1970s~\cite{SmarrThesis,1977NYASA.302..569S}. The field matured in the 1990s, when a large collaboration of research groups worked together toward solving the ``Grand Challenge'' of evolving BBHs~\cite{Abrahams1998}. The crucial final breakthrough came in 2005, when three groups \cite{Pretorius2005a, Campanelli2006a,Baker2006a} devised two completely independent techniques to numerically solve the BBH problem. The first solution \cite{Pretorius2005a} handled the spacetime singularity by excising the regions interior to the black hole horizons, while the second technique \cite{Campanelli2006a,Baker2006a}, dubbed the ``moving punctures approach,'' used singularity avoiding slices of the black hole spacetimes. Since then, through considerable effort by many research groups (reviewed, e.g., in Refs~\cite{Choptuik15,sperhake2014numerical,Tiec:2014lba,Hannam:2013pra,Pfeiffer:2012pc}), each of these techniques have matured, yielding two distinct, independent approaches to modeling BBHs with numerical relativity. Important analytic and numerical developments \cite{Loffler:2011ay,Mroue:2013PRL,Szilagyi:2015rwa} have improved the accuracy and efficiency of the codes implementing each technique. Each technique has enabled numerical relativists to begin constructing catalogs of BBHs and associated gravitational waveforms~\cite{Ajith:2012az,Hinder:2013oqa,Mroue:2013PRL,Pekowsky:2013ska,Healy:2014yta,Clark:2014fva,Husa:2015iqa}. A BBH can be characterized by seven intrinsic parameters: the ratio $q$ of the holes' masses and the spin-angular-momentum vectors $S_1^i$ and $S_2^i$ of each hole; as a result, these catalogs must include many BBH simulations to span the parameter space of BBH events that LIGO could observe. Note that early work has already used numerical-relativity waveforms for detection and parameter estimation in LIGO, for the injection of waveforms~\cite{Aylott:2009ya,Aylott:2009tn,Ajith:2012az,Aasi:2014tra}, into LIGO noise, and for establishing common error measures to establish standard of required waveform accuracy \cite{Hinder:2013oqa}. In this paper, we present a detailed comparison of the targeted numerical BBH simulations that modeled GW150914 in Figs. 1--2 of Ref.~\cite{Abbott:2016blz}, provided by two codes: SpEC and LazEv. We chose the parameters (masses and spins) of these simulations to be consistent with estimates of the parameters for GW150914~\cite{Abbott:2016blz,TheLIGOScientific:2016wfe}. Note that we are not presenting any additional information on parameter estimates in this paper. The parameters we chose are consistent with LIGO's observation of GW150914, but we could have chosen different parameters that are also consistent with the observation. As discussed in Ref.~\cite{Abbott:2016blz,TheLIGOScientific:2016wfe}, there is considerable uncertainty in parameter estimates for GW150914, particularly for the black-hole spins. By comparing the results from these two codes, our study extends the statement made in the caption of Fig.~1 of Ref.~\cite{Abbott:2016blz}: that the numerical relativity calculations shown there are ``confirmed to 99.9\%.'' Our investigation extends previous validation studies of each code internally~\cite{Baker-Campanelli-etal:2007,Hannam:2009hh}, and against one another \cite{2012CQGra..29l4001A}, using common standards for waveform accuracy \cite{Hinder:2013oqa}, to current versions of both numerical-relativity codes for the important case of modeling GW150914. SpEC and LazEv are completely independent. They use different formulations of Einstein's equations, so they assume different decompositions of $G_{\mu\nu}=8\pi T_{\mu\nu}$ into evolution equations and constraints, and they solve for different dynamical variables. They use different methods to choose coordinates and different methods of handling the spacetime singularities inside the black holes. They use different analytic and numerical methods for generating constraint-satisfying initial data, and they use different geometries for the numerical grid. They use different numerical methods for the spacetime evolution, for refining the numerical grid, and for extracting the gravitational waveforms from the evolved variables. They share no subroutines in common. But despite these vast differences, we show in this paper that the two codes produce the same physics. This is a very strong test of both codes and of the analytical and numerical methods underlying them. Moreover, our comparison itself began independently. While we based both simulations on the same mass and spin parameter estimates, we only began discussing the comparison when lower resolutions had finished and higher resolutions were already under way. We made no special effort to tune the codes or the simulations to agree with each other. In this way, we have demonstrated the agreement of our codes under realistic working conditions, where multiple groups independently perform rapid follow up to a LIGO observation. These results build confidence that both numerical methods produce consistent physics, which in turn builds confidence in current and future studies making use of targeted numerical-relativity simulations to follow up LIGO observations. Targeted numerical-relativity simulations already feature in recent studies directly comparing data for GW150914 to catalogs of numerical-relativity simulations~\cite{Abbott:2016apu,Jani:2016wkt} and in recent studies injecting numerical-relativity waveforms into LIGO data for GW150914, to help assess systematic errors in approximate waveform models~\cite{TheLIGOScientific:2016wfe, Abbott:2016izl, TheLIGOScientific:2016src}. This paper is organized as follows. In Sec.~\ref{sec:SXS} we briefly describe the formalism and implementation of SpEC, used by the Simulating eXtreme Spacetimes (SXS) collaboration to numerically evolve BBHs. In Sec.~\ref{sec:RIT} we describe the different formalism and code implementation of LazEv, used by the Rochester Institute of Technology (RIT) group. Table \ref{tab:methods} summarizes the independent methods these two codes employ to construct and evolve initial data for black holes. In Sec.~\ref{sec:results}, we directly compare the waveforms produced by each approach to one another. In Sec.~\ref{sec:discussion} we conclude with a discussion of the significance of those comparisons and implications for future comparisons of observations with numerical-relativity calculations. \begin{table*} \begin{tabular}{|>{\raggedright}p{2.2in}||>{\raggedright}p{2.2in}|>{\raggedright\arraybackslash}p{2.2in}|} \hline & \textbf{LazEv} & \textbf{SpEC} \\ \hline \hline \multicolumn{3}{|l|}{\hspace*{2cm} \bf \emph{Initial data}}\\ \hline Formulation of Einstein constraint equations & conformal method using Bowen-York solutions~\cite{Murchadha-York:1974b,Pfeiffer2003b,Bowen-York:1980}& conformal thin sandwich \cite{York1999,Pfeiffer2003b} \\ \hline Singularity treatment & puncture data~\cite{Brandt1997} & quasi-equilibrium black-hole excision~\cite{Caudill-etal:2006,Cook2004,Cook2002} \\ \hline Numerical method & pseudo-spectral~\cite{AnsorgBruegmann2004} & pseudo-spectral~\cite{Pfeiffer2003} \\ \hline Achieving low orbital eccentricity & post-Newtonian inspiral \cite{Husa-Hannam-etal:2007} & iterative eccentricity removal~\cite{Pfeiffer-Brown-etal:2007,PhysRevD.83.104034} \\ \hline \hline \multicolumn{3}{|l|}{\hspace*{2cm}\bf \emph{Evolution}} \\ \hline Formulation of Einstein evolution equations & BSSNOK~\cite{NOK87,Shibata1995,Baumgarte99} & first-order generalized harmonic with constraint damping~\cite{Lindblom:2007,Friedrich1985,Pretorius2005c,Pretorius2005a} \\ \hline Gauge conditions & evolved lapse and shift~\cite{Bona1997,Alcubierre2002,vanMeter2006} & damped harmonic~\cite{Szilagyi:2009qz} \\ \hline Singularity treatment & moving punctures~\cite{Campanelli2006a,Baker2006a} & excision \cite{Scheel2006} \\ \hline Outer boundary treatment & Sommerfeld & minimally-reflective, constraint-preserving \cite{Lindblom:2007,Rinne2007} \\ \hline Discretization & high-order finite-differences~\cite{Husa2007,Zlochower:2005bj} & pseudo-spectral methods \\ \hline Mesh refinement & adaptive mesh refinement~\cite{Schnetter2003b} & domain decomposition with spectral adaptive mesh refinement~\cite{Pfeiffer2003,Szilagyi:2009qz} \\ \hline \end{tabular} \caption{A comparison of the two independent numerical relativity codes described in the text. Each code uses different techniques to construct and evolve initial data for BBHs and to extract the emitted gravitational radiation. This table is based on Table I of Ref.~\cite{Pfeiffer:2012pc}. \label{tab:methods} } \end{table*}
\label{sec:discussion} We have demonstrated that two completely independent codes to evolve binary black holes (SpEC and LazEv) produce very similar results. As shown in Fig.~\ref{fig:comparison1} and Table \ref{tab:match}, we find good agreement even with moderately low resolution simulations (i.e. N100 and L5). A detailed convergence analysis, like that summarized in Table \ref{tab:match}, suggests that both the generalized harmonic \cite{Pretorius2005a} and moving puncture \cite{Campanelli2006a,Baker2006a} approaches lead to accurate solutions of the general relativity field equations. Given that the initial configurations are not exactly the same (different eccentricities, slightly different masses and spins), we consider this general agreement an excellent verification of the analytic formulations, numerical methods, and code implementations used in both SpEC and LazEv. The next steps in further verifying the results of numerical relativity codes will be to consider binary systems with precession and to consider simulations that follow a larger number of binary orbits. For these more demanding tests, it will be more important to start different codes with closely coordinated initial parameters. This study will be the subject of a future publication. Future work also includes considering cases with more extreme parameters. Here, the simulations' very good agreement with the SEOB waveform is not surprising, since the moderate spins and almost equal masses make this an especially easy region of the parameter space to model. But for higher mass ratios and more extreme spins, numerical relativity might disagree more strongly with semianalytic, approximate waveforms, especially in regions where the semianalytic models have not been tuned to numerical relativity. Recent studies have begun exploring the agreement of numerical relativity and approximate, analytic waveforms in different regions of the BBH parameter space~\cite{Kumar:2016dhh,Jani:2016wkt,Husa:2015iqa,Kumar:2015tha,Szilagyi:2015rwa}. However, from the results of Fig.~\ref{fig:comparisonEOB} we can already conclude that even if analytic waveform models provide a very good approximation to the true prediction of general relativity, full numerical solutions of Einstein's equations can be visibly more accurate than analytic models. Targeted followup with numerical relativity is thus an important tool for comparing gravitational-wave observation and theory and for reliably measuring potential deviations from Einstein's theory of gravitation \cite{TheLIGOScientific:2016src}. Our study suggests that both groups' standard production simulations are sufficiently accurate and efficient to respond rapidly and comprehensively to further events like GW150914, informing the analysis and interpretation of LIGO data. Followup simulations of events like GW150914 can be performed on a timescale of days to weeks (depending on resolution) and at low computational cost, with confidence that both methods produce consistent physics. This is important for the construction of numerically generated waveform data banks with simulations from heterogeneous codes and formalisms. However, numerical-relativity simulations can be considerably more costly and challenging elsewhere in the BBH parameter space, particularly if they remain in LIGO's band for more orbits, such as GW151226, or if they have more extreme parameters. For instance, the SpEC simulation modeling GW151226 that appears in Fig.~5 of Ref.~\cite{Abbott:2016nmj} (SXS:BBH:0317 at \url{http://black-holes.org/waveforms}) required approximately 2 months to complete, and a recent simulation similar to those used here to model GW150914 but with spins $\chi=+0.96$ for the larger black hole and $\chi=-0.9$ for the smaller black hole (SXS:BBH:0306 at \url{http://black-holes.org/waveforms}) required approximately two months to complete. Future work includes enabling more rapid, targeted follow up numerical-relativity simulations for these more challenging cases.
16
7
1607.05377
In fall of 2015, the two LIGO detectors measured the gravitational wave signal GW150914, which originated from a pair of merging black holes (Abbott et al Virgo, LIGO Scientific 2016 Phys. Rev. Lett. <A href="http://journals.aps.org/prl/abstract/10.1103/PhysRevLett.116.061102">116</A> <A href="http://journals.aps.org/prl/abstract/10.1103/PhysRevLett.116.061102">061102</A>). In the final 0.2 s (about 8 gravitational-wave cycles) before the amplitude reached its maximum, the observed signal swept up in amplitude and frequency, from 35 Hz to 150 Hz. The theoretical gravitational-wave signal for merging black holes, as predicted by general relativity, can be computed only by full numerical relativity, because analytic approximations fail near the time of merger. Moreover, the nearly-equal masses, moderate spins, and small number of orbits of GW150914 are especially straightforward and efficient to simulate with modern numerical-relativity codes. In this paper, we report the modeling of GW150914 with numerical-relativity simulations, using black-hole masses and spins consistent with those inferred from LIGO’s measurement (Abbott et al LIGO Scientific Collaboration, Virgo Collaboration 2016 Phys. Rev. Lett. <A href="http://journals.aps.org/prl/abstract/10.1103/PhysRevLett.116.241102">116</A> <A href="http://journals.aps.org/prl/abstract/10.1103/PhysRevLett.116.241102">241102</A>). In particular, we employ two independent numerical-relativity codes that use completely different analytical and numerical methods to model the same merging black holes and to compute the emitted gravitational waveform; we find excellent agreement between the waveforms produced by the two independent codes. These results demonstrate the validity, impact, and potential of current and future studies using rapid-response, targeted numerical-relativity simulations for better understanding gravitational-wave observations.
false
[ "full numerical relativity", "black holes", "Abbott et al LIGO Scientific Collaboration", "al LIGO Scientific Collaboration", "general relativity", "Abbott et al Virgo", "LIGO Scientific", "modern numerical-relativity codes", "al Virgo", "numerical-relativity simulations", "LIGO", "the same merging black holes", "150 Hz", "35 Hz", "merger", "two independent numerical-relativity codes", "black-hole masses", "the gravitational wave signal GW150914", "analytic approximations", "LIGO Scientific 2016 Phys" ]
8.082224
2.308672
62
4056591
[ "Varenius, E.", "Conway, J. E.", "Martí-Vidal, I.", "Aalto, S.", "Barcos-Muñoz, L.", "König, S.", "Pérez-Torres, M. A.", "Deller, A. T.", "Moldón, J.", "Gallagher, J. S.", "Yoast-Hull, T. M.", "Horellou, C.", "Morabito, L. K.", "Alberdi, A.", "Jackson, N.", "Beswick, R.", "Carozzi, T. D.", "Wucknitz, O.", "Ramírez-Olivencia, N." ]
2016A&A...593A..86V
[ "Subarcsecond international LOFAR radio images of Arp 220 at 150 MHz. A kpc-scale star forming disk surrounding nuclei with shocked outflows" ]
49
[ "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Department of Astronomy, University of Virginia, 530 McCormick Road, Charlottesville, VA, 22904, USA", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Instituto de Astrofísica de Andalucía (IAA, CSIC), Glorieta de las Astronomía, s/n, 18008, Granada, Spain; Departamento de Física Teorica, Facultad de Ciencias, Universidad de Zaragoza, Spain", "The Netherlands Institute for Radio Astronomy (ASTRON), PO Box 2, 7990 AA, Dwingeloo, The Netherlands", "Jodrell Bank Centre for Astrophysics, Alan Turing Building, School of Physics and Astronomy, The University of Manchester, Manchester, M13 9PL, UK", "Department of Astronomy, University of Wisconsin-Madison, WI 53706, USA", "Department of Physics, University of Wisconsin-Madison, WI 53706, USA; Center for Magnetic Self-Organization in Laboratory and Astrophysical Plasmas, University of Wisconsin-Madison, WI 53706, USA", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Leiden Observatory, Leiden University, PO Box 9513, 2300 RA, Leiden, The Netherlands", "Instituto de Astrofísica de Andalucía (IAA, CSIC), Glorieta de las Astronomía, s/n, 18008, Granada, Spain", "Jodrell Bank Centre for Astrophysics, Alan Turing Building, School of Physics and Astronomy, The University of Manchester, Manchester, M13 9PL, UK", "Jodrell Bank Centre for Astrophysics, Alan Turing Building, School of Physics and Astronomy, The University of Manchester, Manchester, M13 9PL, UK", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany", "Instituto de Astrofísica de Andalucía (IAA, CSIC), Glorieta de las Astronomía, s/n, 18008, Granada, Spain" ]
[ "2016PhDT.......253B", "2017A&A...602A..42K", "2017A&A...602A..54M", "2017ApJ...841...44P", "2017ApJ...844..108M", "2017ApJ...849...14S", "2017MNRAS.468..946S", "2017MNRAS.469L..89Y", "2017MNRAS.471.1634H", "2018A&A...610L..18R", "2018A&A...613C...3V", "2018A&A...619A..36C", "2018ApJ...853L..28B", "2018arXiv180302831S", "2019A&A...622A...4H", "2019A&A...623A.173V", "2019A&A...627A.147A", "2019ApJ...873...21H", "2019MNRAS.484.3665Y", "2019MNRAS.487..168P", "2019MNRAS.490.5440P", "2020ApJ...888...36R", "2020ApJ...903..138A", "2020MNRAS.493.2817K", "2020MNRAS.496.2613B", "2020PASP..132c5001L", "2021A&ARv..29....2P", "2021ApJ...923..206S", "2021ApJ...923..240S", "2021ApJS..257...42L", "2021MNRAS.500.2620Y", "2021MNRAS.508.5092Y", "2022A&A...657A..49K", "2022A&A...658A...1M", "2022A&A...658A...3S", "2022A&A...658A...4R", "2022A&A...658A...6K", "2022A&A...664A..25K", "2022ApJ...936...58Z", "2022PASJ...74..407U", "2023ApJ...943..142C", "2023ApJ...953..157V", "2023MNRAS.523.5487B", "2023arXiv231113427D", "2024A&A...683A..23D", "2024MNRAS.529.4468L", "2024arXiv240313948P", "2024arXiv240518978R", "2024arXiv240606689G" ]
[ "astronomy" ]
5
[ "ISM: structure", "techniques: high angular resolution", "galaxies: individual: Arp 220", "galaxies: starburst", "Astrophysics - Astrophysics of Galaxies" ]
[ "1966ApJS...14....1A", "1988MNRAS.230..345N", "1988MNRAS.234..919H", "1990ApJS...74..833H", "1990MNRAS.246..110M", "1991ApJ...378...65C", "1991MNRAS.251..112S", "1992ARA&A..30..575C", "1995PhDT.......238B", "1996AJ....111.1945D", "1996MNRAS.282..779W", "1997ApJ...484..702S", "1998ApJ...492L.107S", "1998ApJ...493L..17S", "1998ApJ...498..541K", "1998ApJ...507..615D", "1999ApJ...514...68S", "2000ApJ...537..613A", "2001ApJ...547..792D", "2001ApJ...554..803Y", "2001ApJ...560..160A", "2001ApJ...560..168M", "2003ASSL..285..109G", "2003ApJ...586..794B", "2003ApJ...591..154M", "2003MNRAS.343..585F", "2004ApJ...617..966T", "2005A&A...443..357G", "2006ASPC..351..497K", "2006ApJ...647..185L", "2007A&A...468L..57D", "2007ASPC..376..127M", "2007ApJ...659..314P", "2008ApJ...684..957S", "2009ApJ...700L.104S", "2010ApJ...710.1462W", "2010MNRAS.405..887C", "2011A&A...535A..79H", "2011AJ....141..100H", "2011ApJ...740...95B", "2012A&A...539A..95O", "2012ApJ...754...58K", "2012MNRAS.423L..30S", "2013ApJ...777...58M", "2013MNRAS.430.3171L", "2013MNRAS.431.3003L", "2014ApJ...789L..36W", "2014JKAS...47..195K", "2015A&A...574A..73M", "2015A&A...574A.114V", "2015AJ....149...32M", "2015ApJ...799...10B", "2015ApJ...800...25T", "2015ApJ...800...70S", "2015ApJ...810..149L", "2015MNRAS.453..222Y", "2015MNRAS.453..638D", "2016ApJ...821L..20P", "2016ApJ...823L..17G", "2016MNRAS.457L..29Y" ]
[ "10.1051/0004-6361/201628702", "10.48550/arXiv.1607.02761" ]
1607
1607.02761_arXiv.txt
Arp\,220 is the closest (77\,Mpc) ULIRG and has been extensively studied across the electromagnetic spectrum. It is a late-stage merger, which explains the peculiar morphology noted by \cite{arp1966}. The centre is heavily obscured in the optical, but radio observations reveal two bright sources about 1$''$ (370\,pc) apart, thought to be the nuclei of two merging galaxies \citep{norris1988}. The two nuclei resemble rotating exponential disks with ongoing star formation \citep{sakamoto1999,sakamoto2008,scoville2015,barcosmunoz2015}, giving rise to dozens of supernovae and supernova remnants detected using very long baseline interferometry (VLBI) at cm wavelengths \citep{smith1998,lonsdale2006,parra2007,batejat2011}. While most of the dust and GHz radio continuum comes from the two nuclei \citep{sakamoto2008,barcosmunoz2015}, CO observations reveal extended dense molecular gas surrounding the nuclei in a kpc-scale ring or disk \citep{scoville1997,downes1998,sakamoto2008,koenig2012}. Optical and X-ray observations provide evidence for a bi-conical outflow, or superwind, carrying energy and matter out to several kpc from the centre of the galaxy \citep{heckman1990,arribas2001,mcdowell2003}. While there is plenty of evidence for intense star formation in Arp\,220, the presence of one or more active galactic nuclei (AGN) has not been established. \cite{yun2001} finds Arp\,220 fainter at 1.4\,GHz than expected from the FIR/radio correlation for star forming galaxies. This may be because part of the IR comes from a radio-weak AGN, which could also explain the lack of a clear AGN-candidate amongst the compact objects detected with cm-VLBI. Unfortunately the extreme column densities of $10^{25}$cm$^{-2}$ \citep{wilson2014} towards the nuclei greatly hinders X-ray observations which could offer direct evidence of AGN activity. The lack of radio emission could also be explained if the synchrotron emission from the nuclei is significantly reduced by thermal (free-free) absorption at GHz frequencies. Since this effect is most prominent at lower frequencies (see e.g. \citealt{condon1992}), observations at MHz frequencies may constrain the properties and structure of the absorbing medium. Regardless of the presence of an AGN in Arp\,220, the mechanical energy from the central starburst would likely contribute significantly to the large-scale superwind. This could manifest itself as outflows from the nuclei and the surrounding molecular disk. Indeed, evidence of outflows from the nuclei with speeds of a few hundred km~s$^{-1}$ has been reported \citep{sakamoto2009,tunnard2015}. These outflows could carry synchrotron-emitting cosmic rays (CRs) out from the nuclei where they could potentially be observed at radio frequencies. Star formation in the molecular disk would also give rise to radio emission on kpc-scales, i.e. far outside the nuclei. The intrinsic steep spectrum of unabsorbed synchrotron emission makes a detection of weak extended emission challenging at GHz frequencies, and indeed no detection of kpc-scale radio emission has been reported. Although Arp\,220 has been observed at MHz frequencies before \citep{sopp1991,waldram1996,douglas1996}, none of those studies resolved the galaxy. Subarcsecond resolution is crucial to understand the relative contributions from the emitting and absorbing structures in the centre of Arp\,220. In this paper, we, for the first time, resolve Arp\,220 at 150\,MHz using the international LOFAR telescope. To complement these observations, we combine archival data from the VLA and MERLIN arrays to obtain both high resolution and sensitivity to extended structure at 1.4\,GHz. In Sect. \ref{sect:obs} we describe the data used in this paper. In Sect. \ref{sect:results} we present our results from the observations, in Sect. \ref{sect:modeling} we describe our modelling, and in Sect. \ref{sect:discussion} we discuss the results and modelling in the context of previous studies. Finally, in Sect. \ref{sect:summary}, we summarise our conclusions. Throughout this paper we assume a distance to Arp 220 of 77\,Mpc ($H_0$=70~km~s$^{-1}$\,Mpc$^{-1}$), i.e. $0\farcs1$=37\,pc, and all spectral indices, $\alpha$, are given according to the power law $S_\nu\propto \nu^\alpha$.
\label{sect:discussion} In this section we discuss the results of observations and modelling and compare with literature. \subsection{Low thermal fractions in the nuclei} The thermal fractions at 1\,GHz, required to match the data, are only 0.8\% and 0.4\% for the east and west nuclei respectively. This is an order of magnitude lower than the $9\%\pm3\%$ found by \cite{marvil2015} for a large sample of galaxies, although \cite{murphy2013} finds lower thermal fractions for ongoing mergers. We note that our modelling assumes no foreground thermal absorption of the nuclei, although this could also affect the synchrotron emission from the nuclei and therefore the thermal fractions reported here should be considered upper limits. As noted by \cite{barcosmunoz2015} based on the lack of thermal emission at 33\,GHz, the easiest explanation for the lower-than-expected thermal flux is that a significant fraction of the ionising photons produced by young stars is absorbed by dust. \subsection{An outflow in the eastern nucleus} \label{sect:eastern} CO(2-1) observations of the eastern nucleus suggests it to be a rotating disk seen almost edge on \citep{sakamoto2008,koenig2012}. In Fig.~\ref{fig:co21} we show the structure of CO(2-1) in the nuclei, showing the orientation of the eastern nucleus. The 150\,MHz features detected north and south of the eastern nuclear disk (Sect. \ref{sect:results}) are consistent with an elongated structure where the northern part is viewed through a free-free absorbing medium, possibly the outer part of the nuclear disk, and the southern part is relatively unaffected by this absorbing medium. The fact that the features are on opposite sides of the centre of the nuclear disk indicates that the emission is associated with the disk. \cite{sakamoto2009} presents evidence for outflows in the nuclei with speeds of 100~km~s$^{-1}$ based on P Cygni profiles in HCO$^+$ and CO. We expect that this outflow could carry radio emitting CRs from the nucleus, or accelerate CRs in situ in the outflow due to shocks and turbulence. This would manifest itself as synchrotron emission tracing an elongated or bipolar feature extending outwards from the centre of the disk, similar to what we observe at 150\,MHz. We note, however, that \cite{barcosmunoz2015} fit a P.A. of 54.7$^\circ$ for the eastern disk, and a perpendicular outflow would then have P.A. = -35.3$^\circ$. The elongation we observe implies a P.A.$\sim0^\circ$, i.e. the 150\,MHz feature is not perfectly aligned with the disk. \emph{HST} NICMOS images of Arp\,220 presented by \cite{scoville1998} show a north-south extension of the eastern nucleus at 1.6\,$\mu$m and 2.2\,$\mu$m. The nature of this structure is not clear, but \cite{scoville1998} argue, based on the extinction structure, that the south side is the near side. Similar north-south structure is also seen in the 3.8\,$\mu$m VLT images presented by \cite{gratadour2005}, their Fig. 5. Another galaxy with a bright star forming disk seen almost edge-on is M82. This galaxy was recently found to have bright radio continuum emission at 150\,MHz around its star forming disk \citep{varenius2015}. \cite{varenius2015} interpret this as the base of the outflow seen at larger distances from M82. From their figure 3b (made at 154\,MHz) it is clear that the emission is brighter on the south-east side of M82, presumably because it is closer and less obscured by free-free absorption in the disk which is partly overlapping the north-west emission. The outflow seen in M82 is brightest about 170\,pc (10$''$) from the centre of the star forming disk, i.e. similar to the 110\,pc (0$\farcs3$) offset in the eastern nucleus of Arp\,220 between the peaks at 33\,GHz and 150\,MHz. Based on the consistency with previous evidence for an outflow in the eastern nucleus, the alignment of the 150\,MHz features above and below the disk, and the similarity to the 150\,MHz observations of M\,82, we interpret the 150\,MHz features as evidence for an outflow in the eastern nucleus, with the southern side being the closer (approaching) part. Multiple studies estimate strong magnetic fields strengths of a few mG for the nuclei in Arp\,220 \citep{lacki+beck2013,yoast-hull2016,torres2004}. In a medium with density $n_{\text{H}_2}\approx10^{4}$\,cm$^{-3}$ and magnetic field strength 2\,mG, we expect a cooling time of about 1000\,yrs at GHz frequencies and slightly less at 150\,MHz \citep[their Fig.~1]{lacki+beck2013}. If we assume wind speed of 1\,000~km~s$^{-1}$, i.e. higher than estimated by \cite{sakamoto2009} and \cite{tunnard2015} but less than the 10\,000~km~s$^{-1}$ reached by SNe \citep{batejat2011}, the CRs could travel only 1\,pc before fading. Even if the wind speed was as high as 5\,000~km~s$^{-1}$ the CRs could barely escape the disk (thickness 10\,pc; \citealt{scoville2015}). The fact that we find emission as far as 100\,pc ($0.3\farcs$) south of the eastern nucleus means that the emitting CRs were accelerated tens of parsecs outside the nuclei, due to star formation outside the nuclei and/or shock-acceleration of CRs in the outflow. We note that Arp\,220 has been recently detected in gamma-rays with Fermi \citep{griffin2016,peng2016}. The measured $\gamma$-ray spectrum presented by \cite{peng2016} is roughly an order of magnitude larger than the predictions of \cite{yoast-hull2015}. This supports the conclusion that a significant fraction of the radio emission must come from CRs accelerated outside the nuclei. It is hard to distinguish between a galactic wind plume driven by star formation and structure due to an AGN hidden in the eastern nucleus. Although the eastern nucleus partially similar to M\,82, which is thought to be starburst driven, the non-perpendicular direction of the outflow may be a sign of AGN activity since the M82\, outflow is perpendicular to the disk both in radio and NIR images. However, given that Arp\,220 is an ongoing merger, this discrepancy may be also explained by the interaction forces of the merging process. More data is needed to determine what is powering the outflow in the eastern nucleus. \subsubsection{An outflow in the western nucleus} The velocity structure in the CS molecule \citep{scoville2015} suggests a disk-like rotation for the western nucleus, with a major axis of the disk in the east-west direction. If an outflow carries CRs out perpendicular to the disk, we expect an elongation in the north-south direction, possibly with a shift of the peak of emission at 150\,MHz due to the inclination. This is consistent with the results presented in Sect. \ref{sect:results}, where we see an elongated feature extending $0.9''$ (330\,pc) south of the western nucleus at 150\,MHz, 1.4\,GHz and 6\,GHz. We also find that even though the contribution from extended emission near the western nucleus is significant, the model still underpredicts the brightness of this nucleus by 40\% at 150\,MHz, see Fig.~\ref{fig:nuclei}. This indicates that there is emission close to this nucleus which is not included in the model. A part of this emission likely comes from this elongated component. Although the LOFAR in-band spectral index map (Fig. \ref{fig:spixinband}) shows a clear north-south elongated feature extending about $0.9''$ (330\,pc) from the western nucleus, the spectral index maps from 150\,MHz to 1.4\,GHz and 6\,GHz (Figs. \ref{fig:spix150_14} and \ref{fig:spix150_6}) do not show this feature. However, if there are free electrons in this region (or in the foreground, for example in a surrounding star forming disk), free-free absorption could be important not only for the nuclear disk (as evident from the modelling in Sect. \ref{sect:modeling}) but also for the elongated feature. The turnover frequency where free-free absorption becomes important depends on the surface brightness (e.g. Eq. \ref{eqn:condon}). The elongated feature is much weaker than the western nucleus and consequently would have a lower turnover frequency. In fact, the measured surface brightness of 40$\times 60\mu$\,Jy~beam$^{-1}=14$\,mJy~arcsec$^{-2}$ at 1.4\,GHz would imply a turnover frequency of about 500\,MHz \citep[their Fig 2]{condon1991}. This could explain the positive spectral index detected within the LOFAR band, while also being consistent with the flat or negative spectral index in Figs. \ref{fig:spix150_14} and \ref{fig:spix150_6}, and consequently a very low surface brightness at GHz frequencies. We interpret the elongated feature detected in the western nucleus as evidence of an outflow, inclined so that the southern part is the approaching part of the outflow, consistent with the model suggested by \cite{tunnard2015}, their Fig.~16. We note that our outflow-extension is ten times larger than the 40\,pc reported by \cite{tunnard2015}. This can be explained by shock-acceleration or star formation also outside the western nucleus, as discussed in Sect. \ref{sect:eastern} for the eastern nucleus, as we would not expect CRs accelerated in the western nucleus to keep radiating out to 330\,pc. We note that the \emph{HST} NICMOS images of Arp\,220 presented by \cite{scoville1998} at 1.1\,$\mu$m and 1.6\,$\mu$m also show features north and south of the western nucleus. Similarly to the eastern nucleus, \cite{scoville1998} argue that the southern side is the closer one, based on the extinction. \subsection{Extended star formation and shocks in the superwind} The fact that we see synchrotron emission as far as 1.5\,kpc ($4''$) from the nuclei in Fig.~\ref{fig:LOFAR} means that the emitting CRs must either be produced this far from the centre, or they must have been advected from regions closer to the nuclei by strong winds, possibly driven by the outflows seen in the nuclei. A combination of the two effects is also possible. There is ample evidence for molecular gas several kpc from the nuclei in the form of a disk of radius 0.5-1\,kpc, inclination $~45^\circ$ and P.A. $~45^\circ$ \citep{scoville1997,downes1998,sakamoto1999,koenig2012}. This structure match very well the contours of the extended emission in Fig.~\ref{fig:LOFAR}. On arcsecond scales, the extended emission detected at 150\,MHz matches well the molecular disk, as seen in CO(1-0) Fig.~\ref{fig:co10}, but seem to have little relation to the orientation of the two nuclei. Although outflows from the nuclei may contribute to the extended emission, it is likely that a significant part of it is produced by star formation in the molecular disk itself. While the nuclei are very dense within the central 50\,pc ($n_{\mathrm{H}_2}\sim2\times10^5$\,cm$^{-2}$; \citealt{scoville2015}), the kpc-scale disk is likely less dense and may have weaker magnetic fields \citep{torres2004}. If so, we expect CRs accelerated here to be able to travel further than those accelerated in the nuclei, which could explain the smooth structure of the extended emission. Although CRs may be accelerated in the disk, it is hard to see how they could stay radiating long enough to reach the ``hooks'' of radio emission in Fig.~\ref{fig:LOFAR} at 1.5\,kpc from the centre (which would require a cooling time of $1.5\times10^6$\,yrs assuming a wind speed of 1~000\,km~s$^{-1}$). However, we note that these ``hooks'' appear perpendicular to the kpc-scale disk, and are also roughly aligned with the direction of the outflows from the nuclear disks. We therefore interpret the ``hooks'' as a sign of a kpc-scale outflow, driven both by the nuclear outflows and by star formation in the extended disk. Indeed, Arp\,220 is known to have a large scale bi-conical outflow, or superwind, seen at optical wavelengths and X-rays \citep{heckman1990,arribas2001,mcdowell2003}. For easy comparison with optical data, we show the 150\,MHz emisson as contours plotted on archival HST I-band data (described in Sect. \ref{sect:restobs}) in Fig. \ref{fig:HST}. \begin{figure*}[htbp] \centering \subfigure{ \includegraphics[width=0.4\textwidth]{HSTmanual} \label{fig:HST-RGB} } \subfigure{ \includegraphics[width=0.48\textwidth]{HST-I} \label{fig:HST-I} } \caption{The left panel shows the HST ACS colour image previously shown by \cite{lockhart2015}, where north is up and east is left. Note the tails on the west side and the compressed, almost flat, side of the galaxy towards the north-east, as well as significant extinction by dust towards the centre. The box has sides of approximately $15''$ and indicates the extent of the right panel. The right panel is a zoom showing the F814W (I-band) data with a saturated colour scale for easy comparison to the results presented in \cite{arribas2001}, their Fig.~2. The black contours are the 150\,MHz continuum at the same levels as in Fig.~\ref{fig:LOFAR}. \label{fig:HST} } \end{figure*} Although the CRs carried from the nuclei may have lost most of their energy at 1.5\,kpc from the centre, the mechanical energy output of the nuclear outflows may still be the driving force of the superwind. Thus, the wind (and the magnetic field) would reach much further out than 1.5\,kpc from the centre. Similarly to the nuclear outflows, shocks in this superwind could accelerate CRs to produce emission far out from the centre of the galaxy. We find this a likely explanation of the ``hooks'' in Fig.~\ref{fig:LOFAR}. We note that observations with even higher sensitivity could potentially detect even weaker (and more extended) radio emission. Furthermore, observations with full polarisation calibration could trace the field lines of the outflow similar to the work by \cite{heesen2011}. We note that both ``hooks'' appear bent towards the south-west. Such a bending could be due to movement of the galaxy towards the north-east relative to the surrounding medium. We note that the HST-image, Fig.~\ref{fig:HST-RGB}, shows a flattening towards the north-east, consistent with such motion. \subsubsection{The rate of star formation in the extended disk} In Fig.~\ref{fig:co10} we show the 150\,MHz radio continuum together with the CO(1-0) line emission. Let us assume that the $50\sigma$ contours of CO(1-0) and $5\sigma$ contours at 150\,MHz trace the same extended star forming medium. Using $N_{\mathrm{H}_2}/W_\text{CO}=1.8$ \citep{dame2001} and Eq. 4 of \cite{kennicutt1998} we can derive a star formation rate per area of 10\,M$_\odot\text{\,yr}^{-1}$kpc$^{-2}$ from CO(1-0). By extrapolating the 150\,MHz brightness to 1.4\,GHz ($\alpha=-0.7$) and using Eq. 6 of \cite{bell2003} we derive a value of 2\,M$_\odot\text{\,yr}^{-1}$kpc$^{-2}$. We note that, in addition to all uncertainties involved in applying these equations to the dense and complex environment of Arp\,220, we also expect the CO value to be too high because of blending with the nucleus, and the 150\,MHz value to be too low because of synchrotron losses in the outer parts of the disk. This is consistent with the difference, and although we refrain from any detailed interpretation of these numbers, we note that they are consistent within an order of magnitude and hence do not argue against star formation in the extended disk. \subsection{Estimating the total star-formation rate} \label{sect:SFR} From the flux densities and spectral indices used in the model described in Sect. \ref{sect:modelingres}, we can extrapolate the radio emission from 32.5\,GHz to 1.4\,GHz without the influence of thermal absorption. We obtain 206\,mJy (east), 281\,mJy (west) and 73.2\,mJy (sphere), i.e. a total spectral luminosity of $L_{1.4\text{\,GHz}}=3.99\times10^{23}\text{W~Hz}^{-1}$. Using Eq. 6 by \cite{bell2003} we calculate the star formation rate (SFR) as $5.52\times10^{-22}L_{1.4\text{\,GHz}}\approx220$\,M$_\odot\text{\,yr}^{-1},$ in good agreement with the 240$\pm30$\,M$_\odot\text{\,yr}^{-1}$ calculated by \cite{farrah2003} using the far-infrared luminosity. Scaling the SFR by the flux density of each component at 1.4\,GHz, we find the three components contributing 81\,M$_\odot\text{\,yr}^{-1}$ (east), 110\,M$_\odot\text{\,yr}^{-1}$ (west) and 29\,M$_\odot\text{\,yr}^{-1}$ (sphere). \subsection{Arp\,220 follows the FIR/radio correlation} \label{sect:FIRradio} The mean ratio of far-IR (FIR) emission to radio emission is usually quantified as a logarithmic ratio called the $q$ parameter (e.g. \citealt{yun2001}, their Eq. 5). \cite{yun2001} find Arp\,220 to have $q=2.67$ which is larger than the mean of $2.34$ for the IRAS 2~Jy sample containing over 9000 sources, i.e. showing less radio emission than expected for Arp\,220. However, using the 1.4\,GHz flux density extrapolated in Sect. \ref{sect:SFR} together with IR measurements from the IRAS point source catalogue v2.1\footnote{Obtained via http://irsa.ipac.caltech.edu/Missions/iras.html.} of $S_{60\mu\text{m}}$ = 104.1\,Jy and $S_{100\mu\text{m}}=117.7$\,Jy we obtain (using Eqns. 5 and 6 by \citealt{yun2001}) a value of $q=2.36$ for Arp\,220, closer to the average value. This shows the importance of accounting for thermal absorption even at GHz frequencies. The fact that Arp\,220 follows the FIR/radio correlation indicates that it is powered by star formation rather than AGN activity. This is consistent with the large number of compact supernovae and supernova remnants found with VLBI observations (see e.g. \citealt{batejat2011}). In addition to star formation and AGN activity, tidal shocks in merging systems may heat the dust and gas enough to produce additional FIR and synchrotron emission, thereby possibly affecting the FIR/radio correlation \citep{murphy2013,donevski2015}. \cite{donevski2015} quantify this effect in terms of the $q$-parameter for different classes of mergers, using the same six-stage merger classification scheme as \cite{haan2011} based on HST imaging. We note that Arp\,220 is hard to classify based on optical imaging due to the extreme dust obscuration of the centre. \cite{haan2011} and \cite{murphy2013} classify Arp\,220 differently using the same scheme, as class 6 and 5 respectively. Both these classes are describing a post-merger system with a single nucleus. We know, from radio observations that Arp\,220 has two nuclei, although the separation of 370\,pc is much smaller that the median ULIRG separation of 1.2\,kpc found by \cite{haan2011} and is therefore easily missed in the HST classification scheme. If we re-classify Arp\,220 as class 4, a late ongoing merger with double nuclei and a tidal tail, our measured $q$-value is in good agreement with the $2.31\pm0.17$ expected for this class by \cite{donevski2015}. Hence, by this argument, emission from tidal shocks seem to play a minor role in Arp\,220. Using the international LOFAR telescope we obtain, for the first time, an image of Arp\,220 at 150\,MHz with subarcsecond resolution. We detect emission associated with the two nuclei known from GHz frequencies, but also extended radio emission reaching 1.5\,kpc from the nuclei. The nuclei have positive spectral indices, indicating significant free-free absorption, while the extended emission has an overall spectral index of $-0.7$, typical for optically thin synchrotron radiation. We find that the extended emission accounts for more than 80\% of the total flux density measured at 150\,MHz. We report on elongated features in the two nuclei, extending 0.3'' (110\,pc) from the eastern nuclear disk and $0.9''$ (330\,pc) from the western nuclear disk, and we interpret these features as evidence for outflows. The extended radio emission follows the CO(1-0) distribution and is likely coming from star formation in the kpc-scale molecular disk surrounding the two nuclei. Outflows from this disk, as well as from the nuclei, likely drives the superwind seen in optical and X-rays wavelengths. We find that shock-acceleration of CRs in the outflows, both in the nuclei and in the base of the superwind) is required to explain the extent of the radio emission. We model the Arp\,220 as a three-component model: the nuclei as exponential disks with thermal absorption, surrounded by a uniform sphere of optically thin synchrotron emission. Our model successfully explains the basic shape of the observed integrated spectrum of the galaxy, as well as the spectra of the nuclei, which are well described by a mixed thermal/non-thermal plasma of thermal fractions 0.8\% and 0.4\% at 1 GHz for the east and west nucleus respectively. These values are an order of magnitude lower than the $9\%\pm3\%$ found by \cite{marvil2015} for a large sample of galaxies. Still, our thermal fractions for Arp\,220 are upper limits because we assume no foreground thermal absorption. The low thermal fractions may be explained by dust absorbing a major part of the ionising photons produced by young stars. Our model underpredicts the flux density in the range 200\,MHz to 1\,GHz for Arp\,220, even when including extended emission, indicating that the emission from outflows in the nuclei (which are prominent below 1\,GHz, but are not considered in our simple model) play an important role in this galaxy. When including the extended emission and accounting for absorption effects, we find that Arp\,220 follows the FIR/radio correlation with $q=2.36$, and we estimate a total star formation rate of 220~M$_\odot\text{\,yr}^{-1}$. International LOFAR observations show great promise to detect extended structures such as outflows or radio halos at MHz frequencies, where they are bright due to the synchrotron spectral slope and lack of free-free absorption. Future international LOFAR observations of Arp\,220 using the Low Band Array at 60\,MHz would be very useful to further disentangle the contributions of the different emitting structures in Arp\,220. Dutch-LOFAR observations will in the future be used to detect and study star forming galaxies. Our results show that for LIRGs the emission detected at 150\,MHz does not necessarily come from the main regions of star formation. This implies that unresolved observations of such galaxies at 150\,MHz is unlikely to be useful for deriving star formation rates. Future studies of LIRGs at MHz frequencies would therefore benefit from using international LOFAR baselines to resolve the star forming structure.
16
7
1607.02761
Context. Arp 220 is the prototypical ultra luminous infrared galaxy (ULIRG). Despite extensive studies, the structure at MHz-frequencies has remained unknown because of limits in spatial resolution. <BR /> Aims: This work aims to constrain the flux and shape of radio emission from Arp 220 at MHz frequencies. <BR /> Methods: We analyse new observations with the International Low Frequency Array (LOFAR) telescope, and archival data from the Multi-Element Radio Linked Interferometer Network (MERLIN) and the Karl G. Jansky Very Large Array (VLA). We model the spatially resolved radio spectrum of Arp 220 from 150 MHz to 33 GHz. <BR /> Results: We present an image of Arp 220 at 150 MHz with resolution 0.̋65 × 0.̋35, sensitivity 0.15 mJy beam<SUP>-1</SUP>, and integrated flux density 394 ± 59 mJy. More than 80% of the detected flux comes from extended (6''≈ 2.2 kpc) steep spectrum (α = -0.7) emission, likely from star formation in the molecular disk surrounding the two nuclei. We find elongated features extending 0.3'' (110 pc) and 0.9'' (330 pc) from the eastern and western nucleus respectively, which we interpret as evidence for outflows. The extent of radio emission requires acceleration of cosmic rays far outside the nuclei. We find that a simple three component model can explain most of the observed radio spectrum of the galaxy. When accounting for absorption at 1.4 GHz, Arp 220 follows the FIR/radio correlation with q = 2.36, and we estimate a star formation rate of 220 M<SUB>⊙</SUB> yr<SUP>-1</SUP>. We derive thermal fractions at 1 GHz of less than 1% for the nuclei, which indicates that a major part of the UV-photons are absorbed by dust. <BR /> Conclusions: International LOFAR observations shows great promise to detect steep spectrum outflows and probe regions of thermal absorption. However, in LIRGs the emission detected at 150 MHz does not necessarily come from the main regions of star formation. This implies that high spatial resolution is crucial for accurate estimates of star formation rates for such galaxies at 150 MHz. <P />The reduced images at 150 MHz and 1.4 GHz presented in this paper are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (<A href="http://130.79.128.5">http://130.79.128.5</A>) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/593/A86">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/593/A86</A>
false
[ "star formation rates", "star formation", "steep spectrum outflows", "dust", "radio emission", "spatial resolution", "such galaxies", "high spatial resolution", "probe regions", "thermal absorption", "anonymous ftp", "MHz frequencies", "a star formation rate", "resolution", "outflows", "thermal fractions", "accurate estimates", "Arp", "International LOFAR observations", "the Multi-Element Radio Linked Interferometer Network" ]
12.917967
8.682631
185
12513976
[ "Hindmarsh, Mark", "Rummukainen, Kari", "Weir, David J." ]
2016PhRvL.117y1601H
[ "New Solutions for Non-Abelian Cosmic Strings" ]
23
[ "Department of Physics and Astronomy, University of Sussex, Falmer, Brighton BN1 9QH, United Kingdom; Department of Physics and Helsinki Institute of Physics, P.O. Box 64, FI-00014 University of Helsinki, Finland", "Department of Physics and Helsinki Institute of Physics, P.O. Box 64, FI-00014 University of Helsinki, Finland", "Institute of Mathematics and Natural Sciences, University of Stavanger, 4036 Stavanger, Norway" ]
[ "2016PhRvD..94d5015W", "2017PhRvD..95f3520H", "2017PhRvD..96b3525H", "2017PhRvD..96j3535R", "2017PhRvD..96l3512A", "2018EL....12261001B", "2018PhLB..778...22B", "2018PhLB..780..485B", "2018PhRvD..98f3519R", "2018PhRvD..98j3533H", "2018arXiv180605814X", "2019FrP.....7...76C", "2019RPPh...82g6901M", "2020EPJC...80..352C", "2020EPJC...80..576A", "2020PhRvD.102h3516R", "2020PhRvD.102j3516F", "2021EPJP..136..618B", "2021NatSR..11.6462X", "2022IJMPD..3150080M", "2023PhRvD.108e5022F", "2023PhRvD.108i5014B", "2024JPCM...36y5602S" ]
[ "astronomy", "physics" ]
16
[ "High Energy Physics - Theory", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Phenomenology" ]
[ "1973NuPhB..61...45N", "1974JETPL..20..194P", "1974NuPhB..79..276T", "1976JPhA....9.1387K", "1982PhLB..113..237K", "1985NuPhB.249..557W", "1985PhLB..153..243W", "1985PhRvL..55.2398H", "1986PhRvD..34.3206D", "1986PhRvL..56.2564D", "1987PhRvD..35.3105A", "1987PhRvL..58.1910C", "1989NuPhB.323..209D", "1989PhRvL..62.1948L", "1990PhRvL..64.1632A", "1990PhRvL..65..668A", "1991PhRvD..44.3067V", "1992PhyB..178...47H", "1993NuPhB.392..461H", "1995RPPh...58..477H", "1997PhRvL..79.5202B", "1998PhRvL..80.2277V", "2001NuPhB.595..402S", "2001PhRvD..65b3503M", "2002PhLB..536..185S", "2003PhRvD..68j3514J", "2004JHEP...06..013C", "2005PhRvD..71c5002F", "2006JHEP...08..066H", "2007PhRvD..75f5015B", "2008JHEP...02..037U", "2010JCAP...05..014B", "2010PhRvD..82f5004B", "2011CQGra..28t4009C", "2011PThPS.190..197H", "2015JPhG...42i4002K", "2016JCAP...04..009A", "2016JCAP...04..053L", "2016PhRvD..93j5021A", "2017PhRvD..95f3520H" ]
[ "10.1103/PhysRevLett.117.251601", "10.48550/arXiv.1607.00764" ]
1607
1607.00764_arXiv.txt
16
7
1607.00764
We study the properties of classical vortex solutions in a non-Abelian gauge theory. A system of two adjoint Higgs fields breaks the SU(2) gauge symmetry to Z<SUB>2</SUB> , producing 't Hooft-Polyakov monopoles trapped on cosmic strings, termed beads; there are two charges of monopole and two degenerate string solutions. The strings break an accidental discrete Z<SUB>2</SUB> symmetry of the theory, explaining the degeneracy of the ground state. Further symmetries of the model, not previously appreciated, emerge when the masses of the two adjoint Higgs fields are degenerate. The breaking of the enlarged discrete symmetry gives rise to additional string solutions and splits the monopoles into four types of "semipole": kink solutions that interpolate between the string solutions, classified by a complex gauge-invariant magnetic flux and a Z<SUB>4</SUB> charge. At special values of the Higgs self-couplings, the accidental symmetry broken by the string is continuous, giving rise to supercurrents on the strings. The SU(2) theory can be embedded in a wide class of grand unified theories (GUTs), including SO(10). We argue that semipoles and supercurrents are generic on GUT strings.
false
[ "additional string solutions", "cosmic strings", "GUT strings", "kink solutions", "classical vortex solutions", "two degenerate string solutions", "monopole", "grand unified theories", "the string solutions", "termed beads", "a non-Abelian gauge theory", "non-Abelian", "SO(10", "t Hooft-Polyakov monopoles", "The strings", "the string", "the strings", "Further symmetries", "supercurrents", "a complex gauge-invariant magnetic flux" ]
10.219341
-1.401841
86
12472810
[ "Petts, J. A.", "Read, J. I.", "Gualandris, A." ]
2016MNRAS.463..858P
[ "A semi-analytic dynamical friction model for cored galaxies" ]
58
[ "Department of Physics, University of Surrey, Guildford GU2 7XH, UK", "Department of Physics, University of Surrey, Guildford GU2 7XH, UK", "Department of Physics, University of Surrey, Guildford GU2 7XH, UK" ]
[ "2017ApJ...837...54C", "2017ApJ...840...31D", "2017ApJ...844...64A", "2017MNRAS.467.3775P", "2017MNRAS.467.4491I", "2017MNRAS.469.2845D", "2017MNRAS.471..478A", "2017PhRvD..96f3001G", "2018A&A...615A..91B", "2018ApJ...864L..19T", "2018ApJ...867..119F", "2018ApJ...868..134K", "2018MNRAS.476.3124C", "2018NewAR..81....1R", "2018PhR...761....1B", "2019ApJ...877..133D", "2019ApJ...881..106M", "2019MNRAS.483..152A", "2019MNRAS.483.4724F", "2019MNRAS.488.2977O", "2019NewAR..8601525D", "2020ApJ...903..149D", "2020MNRAS.492.5102L", "2020MNRAS.493..320L", "2020MNRAS.494.3053B", "2020MNRAS.498.3601B", "2020MNRAS.499..804G", "2021ApJ...912...43B", "2021ApJ...916....9M", "2021ApJ...916...55T", "2021ApJ...919..140S", "2021Galax..10....5B", "2021MNRAS.501..179M", "2021MNRAS.502.3554B", "2021MNRAS.507.4997E", "2022ApJ...926..215B", "2022JPhG...49f3001M", "2022MNRAS.510.2900D", "2022MNRAS.511.1860S", "2022MNRAS.511.2339Q", "2022MNRAS.511.4753G", "2022MNRAS.512..739D", "2022MNRAS.512.3365B", "2022MNRAS.515..185O", "2022MNRAS.515..407K", "2022arXiv220710638A", "2023A&A...680A..91K", "2023ApJ...950..178M", "2023MNRAS.519..948L", "2023MNRAS.522..948C", "2023MNRAS.523.2721H", "2023PhRvD.108j3011D", "2024A&A...681A..73T", "2024AJ....167...76L", "2024JCAP...02..054G", "2024arXiv240416184A", "2024arXiv240508870G", "2024arXiv240518468O" ]
[ "astronomy" ]
9
[ "methods: numerical", "galaxies: kinematics and dynamics", "Astrophysics - Astrophysics of Galaxies" ]
[ "1943ApJ....97..255C", "1959AnAp...22..126H", "1960AnAp...23..474H", "1962AJ.....67..471K", "1983ApJ...274...53W", "1987MNRAS.224..349B", "1988MNRAS.235..289Z", "1993MNRAS.262.1062S", "1993MNRAS.265..250D", "1994AJ....107..634T", "1997MNRAS.289..253C", "2003ApJ...582..196H", "2005A&A...431..861J", "2005MNRAS.364.1105S", "2006MNRAS.368.1073G", "2006MNRAS.373.1451R", "2008ApJ...678..780G", "2008gady.book.....B", "2009MNRAS.397..709I", "2010ApJ...725.1707G", "2010MNRAS.405.2327P", "2010RAA....10.1242G", "2011MNRAS.411..653J", "2011MNRAS.416.1181I", "2012ApJ...745...83A", "2012MNRAS.426..601C", "2014ApJ...785...51A", "2014MNRAS.444.3738A", "2015ApJ...806..220A", "2015MNRAS.454.3778P", "2016JCAP...05..021S", "2016MNRAS.457.1339P" ]
[ "10.1093/mnras/stw2011", "10.48550/arXiv.1607.04284" ]
1607
1607.04284_arXiv.txt
Dynamical friction is a drag force caused by momentum exchange between a massive object moving within a sea of lighter background `stars'. (We shall refer to these as `stars' throughout this paper, though they could comprise any self-gravitating entity; e.g. dark matter particles \citep{BT08}.) Dynamical friction is thought to be a key driver of the mergers of star clusters, galaxies, and even galaxy clusters over cosmic time \citep{Gan10,Peirani10,ArcaSedda14c,ArcaSedda15,Priyatikanto16}. In a seminal work \citet{Chandrasekhar43} calculated the force on a massive object traversing an infinite homogeneous isotropic background. Approximating the drag to come only from stars moving slower than the satellite and ignoring the velocity dependence of the strength of interactions, Chandrasekhar's formula is often seen in the following simplified analytic form \citep{Chandrasekhar43, BT08}: \begin{equation} \frac{d\bmath{v\sscript{S}}}{dt} = -4\pi G^2 M\sscript{S} \rho \log(\Lambda)f(v\sscript{*}<v\sscript{s})\frac{\bmath{v\sscript{S}}}{v\sscript{S}^3}, \label{dynfric.eq} \end{equation} where $\log(\Lambda)$ is the Coulomb logarithm equal to $ \log{(b\sscript{max}/b\sscript{min})}$ where $b\sscript{max}$ and $b\sscript{min}$ are the maximum and minimum impact parameters for encounters, and it is assumed that $b\sscript{max} \gg b\sscript{min}$. $M\sscript{S}$ and $\bmath{v\sscript{S}}$ are the satellite's mass and velocity ($v\sscript{S} = \bmath{|v\sscript{S}|}$), $\rho$ is the background density and $f(v\sscript{*}<v\sscript{s})$ is the fraction of stars moving slower than the satellite. Although derived under very simple assumptions, Chandrasekhar's formula has proven to be remarkably successful for more general isotropic spherical distributions. Chandrasekhar's formula has been shown to be a good approximation for low mass satellites due to the fact that the global response from the background appears to be negligible \citep{White83,Bontekoe87,Zaritsky88,Cora97}. For satellites more massive than $\gtrsim 10\%$ of the host galaxy mass, the global response is non-negligible. As such equation \ref{dynfric.eq} is not accurate for major mergers. However, for large mass ratios (e.g. dwarf galaxies moving in the halo of larger hosts, globular clusters in dwarf spheroidals, etc.), the formula has been very successful at reproducing the orbital evolution. Although the formula has remained largely unchanged, its application has seen many improvements in recent years \citep{Hashimoto03,Just05,Just11,Petts15}. Typically, $\log(\Lambda)$ is poorly defined, with $b\sscript{max}$ being of the order of the size of the system and $b\sscript{min}$ the impact parameter for a 90 degree deflection \citep{BT08}. Because the impact parameters are found inside a logarithm which is slowly varying (so long as $b\sscript{max} \gg b\sscript{min}$), a constant $\log(\Lambda)$ is usually assumed, such that: \begin{equation} \log(\Lambda) \sim \log\left(\frac{r\sscript{galaxy}}{\max(r\sscript{hm},\frac{GM\sscript{s}}{v^2\sscript{typ}})}\right), \end{equation} where $r\sscript{galaxy}$ is the effective ``size'' of the galaxy, and the minimum impact parameter is the larger of the size of the satellite and the impact parameter for which a star interacting at a ``typical velocity'', $v\sscript{typ}$, is deflected by 90 degrees \citep{BT08}. \citet{Just11} showed that the approximation of local homogeneity is satisfied when $b\sscript{max}$ is set to be the length scale over which the density can be assumed to be constant \citep[see also][]{Just05}: \begin{equation} b\sscript{max}(r) = \rm{min}\left(\frac{\rho(r)}{|\nabla\rho|},r\right), \end{equation} where $r$ is the galactocentric distance to the satellite, i.e. $b\sscript{max}$ varies with galactocentric distance and the slope of the profile \citep{Just05,Just11}. However, this prescription neglects the force from particles deep in the cuspy region of the galaxy, and under-predicts the drag when the satellite is very close to the centre. \citep[See][for a method for dealing with steep galactic centres.]{ArcaSedda14} A radially varying $\log(\Lambda)$ improves the agreement with $N-$body models of the inspiral of eccentric satellites, where the satellite experiences a large variation in its radial position over an orbital time \citep[e.g.][]{Hashimoto03}. The variations of $\log(\Lambda)$ are also especially important when $b\sscript{max}/b\sscript{min}$ approaches unity \citep{Petts15}, which for most systems occurs at approximately $M\sscript{s} \sim M\sscript{enc}$ where the satellite is assumed to form a quasi two-body system with the galactic centre and dynamical friction becomes inefficient \citep[e.g.][]{BT08, Gualandris08}. For galaxies with a large core, the inspiral of satellites stalls further out, at the edge of the constant density core \citep{Read06, Goerdt06,Inoue09,Goerdt10,Cole12}. In \citet{Petts15} we introduced a semi-analytic model that reproduces the inspiral and correct stalling radii of satellites orbiting a Dehnen background \citep{Dehnen93}, including the case where the asymptotic logarithmic slope approached zero. However, the model fared less well for large constant density cores. It also failed to reproduce the rapid ``super-Chandrasekhar'' phase reported in \citet{Read06} (hereafter R06). In this paper we show that this dramatic ``core-stalling'' effect can be approximately captured if we consider the radius at which the satellite tidally disrupts the core and sculpts the velocity distribution in this region. This idea was already explored in \citet{Goerdt10} for massive satellites infalling within relatively cuspy background distributions. Here, we show that this same idea can be generalised to large constant density cores in which the tidal radius of the satellite approaches the size of the cored region. We also address the ``Super-Chandrasekhar'' friction phase observed in R06 previously thought to be ``super-resonance'' of the harmonic core. The idea was that as the angular frequency in a perfectly flat core is the same for every star, some global resonant effects may drive the friction in a way that cannot be correctly described by considering only two body interactions (R06). As real systems are never truly harmonic -- especially if one considers the back reaction of the satellite -- we argue here that the friction cannot owe to ``super-resonance'' (i.e. a proposed efficient resonant interaction that dominates the friction, see R06). Instead, we show that this phase of rapid infall is due to previously invalid assumptions about the velocity distribution in the core, and can be explained entirely through local friction via two-body interactions. The paper is organised as follows. In section \S \ref{models.ch} we describe the galaxy models used in this study. In section \S \ref{theory.ch} we explain the theory and necessary improvements to our model. In section \S \ref{simulations.ch} we describe the simulations used to test our model. In section \S \ref{results.ch} we compare the results of our new model to $N-$body results. In section \S \ref{discussion.ch} we discuss the stalling mechanism and the potentially related problem of ``dynamical buoyancy'' reported recently in \citet{Cole12}. Finally, in \S \ref{conclusion.ch}, we present our conclusions.
\label{conclusion.ch} We have shown that Chandrasekhar's dynamical friction model considering only two-body encounters is sufficient to explain the inspiral of satellites in constant density profiles, so long as one uses the self-consistent distribution function of velocities instead of the usually assumed Maxwellian distribution. In particular, we show that we can reproduce the ``super-Chandrasekhar'' phase, suggesting that it does not owe to resonance. The Chandrasekhar formula probably works so well because the potential is never truly harmonic in any physically reasonable distribution. The agreement is improved further by including the usually neglected contribution of the fast moving stars, which contributes a non-negligible portion of the drag inside the core. However, even after including the correct background distribution function and the effect of fast moving stars, we find that we are not able to reproduce the stalling observed in large constant density cores such as in the Hénon Isochrone Model. Following \citet{Goerdt10}, we show that such large-core stalling occurs in the same manner as it does for cusps. The infalling satellite tidally disrupts the core when $r\sscript{t}(r\sscript{a}) = r\sscript{a}$. For cusped background models this occurs at $M\sscript{s} \sim M\sscript{enc}$. However, for cored backgrounds, the satellite tidal radius can become very large indeed. This leads to stalling at many times the radius at which the mass in background stars approaches the satellite mass. In our model, $f(v\sscript{*})$ is derived from the distribution function of the background density distribution and our model has no free or tuned parameters. As such, it should be general for any model with a cored or cusped centre. Finally, we suggest that the dynamical friction core-stalling can be understood in two different ways. For a perfectly harmonic background with a single point mass satellite, R06 demonstrated that there exist stable solutions with no net momentum exchange between the satellite and the background. We generalised this model in section \ref{discussion.ch}, showing that the same should be true when multiple satellites are present. While the satellite and the background likely reach an approximation to this state, the correspondence cannot be perfect. Secondly, the ``dynamical buoyancy'' effect discussed in \citet{Cole12} is not captured by our semi-analytic model, nor by the R06 stalling state. Instead, the answer may lie in the frictional force coming from stars moving faster than the satellite. \citet{Inoue11} showed that strongly interacting ``Horn'' stars can both decelerate (P-horn) and accelerate (N-horn) the satellite depending on their relative orbital phase. For a large-cored background, the cumulative effects of P and N-horn stars approximately cancel the friction experienced in the core region, leading to core-stalling as in the R06 model. However, configurations can exist where it is possible for N-horn stars to dominate over P-horn ones, if a satellite begins its life deep inside the constant density core. This is a possible explanation for \citet{Cole12}'s dynamical buoyancy; however, a full proof will require further investigation beyond the scope of this work.
16
7
1607.04284
We present a dynamical friction model based on Chandrasekhar's formula that reproduces the fast inspiral and stalling experienced by satellites orbiting galaxies with a large constant density core. We show that the fast inspiral phase does not owe to resonance. Rather, it owes to the background velocity distribution function for the constant density core being dissimilar from the usually assumed Maxwellian distribution. Using the correct background velocity distribution function and our semi-analytic model from previous work, we are able to correctly reproduce the infall rate in both cored and cusped potentials. However, in the case of large cores, our model is no longer able to correctly capture core-stalling. We show that this stalling owes to the tidal radius of the satellite approaching the size of the core. By switching off dynamical friction when r<SUB>t</SUB>(r) = r (where r<SUB>t</SUB> is the tidal radius at the satellite's position), we arrive at a model which reproduces the N-body results remarkably well. Since the tidal radius can be very large for constant density background distributions, our model recovers the result that stalling can occur for M<SUB>s</SUB>/M<SUB>enc</SUB> ≪ 1, where M<SUB>s</SUB> and M<SUB>enc</SUB> are the mass of the satellite and the enclosed galaxy mass, respectively. Finally, we include the contribution to dynamical friction that comes from stars moving faster than the satellite. This next-to-leading order effect becomes the dominant driver of inspiral near the core region, prior to stalling.
false
[ "constant density background distributions", "large cores", "SUB", "satellites", "a large constant density core", "stalling", "the constant density core", "galaxies", "the correct background velocity distribution function", "the background velocity distribution function", "core-stalling", "previous work", "our semi-analytic model", "Maxwellian", "the core region", "M<SUB>s</SUB>/M", "r", "the enclosed galaxy mass", "a dynamical friction model", "s</SUB" ]
9.596414
5.634005
-1
1118259
[ "Dent, James B.", "Dutta, Bhaskar", "Newstead, Jayden L.", "Strigari, Louis E." ]
2017PhRvD..95e1701D
[ "Dark matter, light mediators, and the neutrino floor" ]
34
[ "Department of Physics, University of Louisiana at Lafayette, Lafayette, Louisiana 70504, USA; Kavli Institute for Theoretical Physics, University of California, Santa Barbara, California 93106-4030, USA", "Mitchell Institute for Fundamental Physics and Astronomy, Department of Physics and Astronomy, Texas A&amp;M University, College Station, Texas 77845, USA", "Department of Physics, Arizona State University, Tempe, Arizona 85287, USA", "Mitchell Institute for Fundamental Physics and Astronomy, Department of Physics and Astronomy, Texas A&amp;M University, College Station, Texas 77845, USA" ]
[ "2017JHEP...04..073B", "2017PhRvD..95c5017F", "2017PhRvD..96i5007D", "2017arXiv171206675P", "2018JCAP...07..009G", "2018JHEP...11..203I", "2018PhLB..785..610M", "2018PhRvD..97c5009D", "2018PhRvD..98j3006O", "2018arXiv180408995P", "2018arXiv180411319P", "2019ARNPS..69..137D", "2019JCAP...01..043B", "2019JHEP...12..124S", "2019PhR...797....1E", "2019arXiv190700991B", "2019arXiv191012437A", "2020PhRvD.101a5012B", "2020PhRvD.102f3024O", "2021PhRvD.104f3013C", "2021PhRvD.104k5022N", "2021PhRvL.127y1802O", "2022LNP...996.....D", "2022PhLB..82937050C", "2022arXiv220305914O", "2022arXiv220307361A", "2022arXiv221109669K", "2023JPhG...50a3001A", "2023PDU....4101242B", "2023ScPPL..71.....F", "2024PhRvD.109d3055Z", "2024PhRvD.109h3016C", "2024arXiv240218454Z", "2024arXiv240403690L" ]
[ "astronomy", "physics" ]
4
[ "High Energy Physics - Phenomenology", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "2010JCAP...11..042F", "2011EPJC...71.1554C", "2011PhRvL.107m1302A", "2013EPJC...73.2501C", "2013JCAP...02..004F", "2013PhRvD..88g6011N", "2013arXiv1311.0029E", "2014PDU.....5...45P", "2014PhRvC..89f5501A", "2014PhRvD..89b3524B", "2014PhRvD..89l3521G", "2014PhRvD..90h3510R", "2014PhRvL.113b1302K", "2015JCAP...03..032L", "2015JCAP...07..026C", "2015JCAP...10..055D", "2015JCAP...12..057G", "2015PhRvD..92f3515D", "2015arXiv150902910T", "2015arXiv151000999M", "2016JCAP...04..027A", "2016JCAP...11..017A", "2016JHEP...05..118C", "2016PhRvD..93g5018D", "2016PhRvD..93j3534S", "2016PhRvL.116g1301A", "2016PhRvL.116p1301A", "2016PhRvL.117g1803F", "2016arXiv160808632A", "2017PhRvD..95c5017F" ]
[ "10.1103/PhysRevD.95.051701", "10.48550/arXiv.1607.01468" ]
1607
1607.01468_arXiv.txt
16
7
1607.01468
We analyze direct dark matter detection experiments for ≲100 MeV mass mediators with general interactions. We compare the nuclear recoil energy spectra from these interactions to the solar neutrino spectrum. A set of interactions that generate spectra similar to the neutrino background is identified; however, this set is distinct from those that mimic the neutrino background for ≳100 MeV mass mediators. We outline a classification scheme based on momentum dependence of the dark matter-nucleus interaction to determine how strong the discovery limit for each interaction saturates due to the neutrino background. Our results motivate experimental progress towards lower nuclear recoil energy thresholds.
false
[ "general interactions", "interactions", "≳100 MeV mass mediators", "≲100 MeV mass mediators", "MeV", "lower nuclear recoil energy thresholds", "the neutrino background", "each interaction saturates", "direct dark matter detection experiments", "these interactions", "the dark matter-nucleus interaction", "spectra", "momentum dependence", "the nuclear recoil energy spectra", "experimental progress", "the solar neutrino", "how strong the discovery limit", "A set", "this set", "a classification scheme" ]
7.600669
-1.853043
42
12593450
[ "Sánchez, Ariel G.", "Scoccimarro, Román", "Crocce, Martín", "Grieb, Jan Niklas", "Salazar-Albornoz, Salvador", "Dalla Vecchia, Claudio", "Lippich, Martha", "Beutler, Florian", "Brownstein, Joel R.", "Chuang, Chia-Hsun", "Eisenstein, Daniel J.", "Kitaura, Francisco-Shu", "Olmstead, Matthew D.", "Percival, Will J.", "Prada, Francisco", "Rodríguez-Torres, Sergio", "Ross, Ashley J.", "Samushia, Lado", "Seo, Hee-Jong", "Tinker, Jeremy", "Tojeiro, Rita", "Vargas-Magaña, Mariana", "Wang, Yuting", "Zhao, Gong-Bo" ]
2017MNRAS.464.1640S
[ "The clustering of galaxies in the completed SDSS-III Baryon Oscillation Spectroscopic Survey: Cosmological implications of the configuration-space clustering wedges" ]
165
[ "Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany", "Center for Cosmology and Particle Physics, Department of Physics, New York University, New York, NY 10003, USA", "Institut de Ciències de l'Espai, IEEC-CSIC, Campus UAB, Carrer de Can Magrans, s/n, E-08193 Bellaterra, Barcelona, Spain", "Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany; Universitäts-Sternwarte München, Ludwig-Maximilians-Universität München, Scheinerstrasse 1, D-81679 Munich, Germany", "Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany; Universitäts-Sternwarte München, Ludwig-Maximilians-Universität München, Scheinerstrasse 1, D-81679 Munich, Germany", "Instituto de Astrofísica de Canarias, C/ Vía Láctea s/n, E-38205 La Laguna, Tenerife, Spain; Departamento de Astrofísica, Universidad de La Laguna, Av. del Astrofísico Francisco Sánchez s/n, E-38206 La Laguna, Tenerife, Spain", "Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany; Universitäts-Sternwarte München, Ludwig-Maximilians-Universität München, Scheinerstrasse 1, D-81679 Munich, Germany", "Institute of Cosmology and Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK; Lawrence Berkeley National Laboratory, 1 Cyclotron Road, Berkeley, CA 94720, USA", "Department of Physics and Astronomy, University of Utah, 115 S 1400 E, Salt Lake City, UT 84112, USA", "Instituto de Física Teórica, UAM/CSIC, Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Leibniz-Institut für Astrophysik Potsdam, An der Sternwarte 16, D-14482 Potsdam, Germany", "Harvard-Smithsonian Center for Astrophysics, 60 Garden St, Cambridge, MA 02138, USA", "Lawrence Berkeley National Laboratory, 1 Cyclotron Road, Berkeley, CA 94720, USA; Leibniz-Institut für Astrophysik Potsdam, An der Sternwarte 16, D-14482 Potsdam, Germany; Departments of Physics and Astronomy, University of California, Berkeley, CA 94720, USA", "Department of Chemistry and Physics, King's College, 133 North River St, Wilkes Barre, PA 18711, USA", "Institute of Cosmology and Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK", "Instituto de Física Teórica, UAM/CSIC, Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Campus of International Excellence UAM+CSIC, Cantoblanco, E-28049 Madrid, Spain; Instituto de Astrofísica de Andalucía, CSIC, Glorieta de la Astronomía, E-18080 Granada, Spain", "Instituto de Física Teórica, UAM/CSIC, Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Campus of International Excellence UAM+CSIC, Cantoblanco, E-28049 Madrid, Spain; Departamento de Física Teórica, Universidad Autónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain", "Institute of Cosmology and Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK; Center for Cosmology and Astro-Particle Physics, Ohio State University, Columbus, OH 43210, USA", "Institute of Cosmology and Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK; Kansas State University, Manhattan, KS 66506, USA; National Abastumani Astrophysical Observatory, Ilia State University, 2A Kazbegi Ave, GE-1060 Tbilisi, Georgia", "Department of Physics and Astronomy, Ohio University, 251B Clippinger Labs, Athens, OH 45701, USA", "Center for Cosmology and Particle Physics, Department of Physics, New York University, New York, NY 10003, USA", "School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews KY16 9SS, UK", "Instituto de Física, Universidad Nacional Autónoma de México, Apdo. Postal 20-364, 01000 México, D.F., Mexico; Department of Physics, Carnegie Mellon University, 5000 Forbes Ave, Pittsburgh, PA 15217, USA; McWilliams Center for Cosmology, Carnegie Mellon University, 5000 Forbes Ave, Pittsburgh, PA 15217, USA", "Institute of Cosmology and Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK; National Astronomy Observatories, Chinese Academy of Science, Beijing 100012, People's Republic of China", "Institute of Cosmology and Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK; National Astronomy Observatories, Chinese Academy of Science, Beijing 100012, People's Republic of China" ]
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[ "astronomy" ]
8
[ "cosmological parameters", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stw2443", "10.48550/arXiv.1607.03147" ]
1607
1607.03147_arXiv.txt
\label{sec:intro} Measurements of the large-scale clustering of galaxies offer a powerful route to obtain accurate cosmological information \citep{Davis1983,Maddox1990,Tegmark2004,Cole2005,Eisenstein2005,Anderson2012,Anderson2013,Anderson2014}. Two-point statistics such as the power spectrum, $P(k)$, and its Fourier transform, the two-point correlation function $\xi(s)$, have been the preferred tools for analyses of the large-scale structure (LSS) of the Universe. The shape of these measurements can be used to constrain the values of several cosmological parameters, providing clues about the nature of dark energy, potential deviations from the predictions of general relativity (GR), the physics of inflation, neutrino masses, etc. \citep{Percival2002,Percival2010,Tegmark2004,Sanchez2006,Sanchez2009,Sanchez2012,Blake2011, Parkinson2012}. A particularly important source of cosmological information contained in the large-scale galaxy clustering pattern is the signature of the baryon acoustic oscillations (BAO), which are the vestige of acoustic waves that propagated through the photon-baryon fluid prior to recombination. The BAO signature was first detected by \citet{Eisenstein2005} in the correlation function of the luminous red galaxy sample of the Sloan Digital Sky Survey \citep[SDSS,][]{York2000}, where it can be seen as a broad peak on large scales \citep{Matsubara2004}, and by \citet{Cole2005} in the power spectrum of the Two-degree Field Galaxy Redshift survey \citep[2dFGRS,][]{Colless2001,Colless2003}, where it appears as a series of wiggles \citep{Eisenstein1998,Meiksin1999}. The position of the peak in the correlation function and the wavelength of the oscillations in the power spectrum closely match the sound horizon scale at the drag redshift, $r_{\rm d}\simeq150\,{\rm Mpc}$. This means that the BAO scale inferred from the clustering of galaxies in the directions parallel and perpendicular to the line of sight can be used as a standard ruler to measure the Hubble parameter, $H(z)$, and the angular diameter distance, $D_{\rm M}(z)$, through the Alcock--Paczynski (AP) test \citep{Alcock1979,Blake2003,Linder2003}. As the AP test cannot be applied to angle-averaged clustering measurements, the full power of the BAO signal can only be exploited by means of anisotropic clustering measurements. That means measurements of the full two-dimensional correlation function or power spectrum \citep{Wagner2008,Shoji2009}, their Legendre multipole moments \citep{Padmanabhan2008} or the clustering wedges statistic \citep*{Kazin2012}. These measurements are affected by redshift-space distortions (RSD) due to the peculiar velocities of the galaxies along the line of sight, which are significantly larger than the geometric distortions due to the AP effect and must be accurately modelled to avoid introducing systematic errors in the obtained constraints. However, more than a complication for the application of the AP test, RSD provide additional cosmological information, as they can be used to constrain the growth rate of cosmic structures \citep{Guzzo2008}. In this way, thanks to the joint information from BAO and RSD, anisotropic clustering measurements can provide information on the expansion history of the Universe and the growth rate of density fluctuations, which is essential to distinguish between dark energy and modified gravity as the driver of cosmic acceleration. Previous analyses of anisotropic clustering measurements based on data from the SDSS-III \citep{Eisenstein2011} Baryon Oscillation Spectroscopic Survey \citep[BOSS,][]{Dawson2013}, clearly illustrated their constraining power \citep{Anderson2013,Anderson2014,Reid2012,Samushia2013,Samushia2014,Chuang2013a, Beutler2014}. In particular, \citet{Sanchez2013,Sanchez2014} explored the cosmological implications of the full shape of measurements of two clustering wedges based on the galaxy samples of BOSS DR11. In this paper we extend these analyses to the final galaxy samples from BOSS, corresponding to SDSS DR12 \citep{Alam2015}. The volume probed by DR12 is only $\sim10$ per cent larger than that of DR11. For this reason, we focus on improving our analysis methodology in order to maximize the cosmological information extracted from the sample. We make use of the joint information of the LOWZ and CMASS galaxy samples into the {\it combined} BOSS sample described in \citet{Reid2016}, increasing the effective volume of the survey with respect to the separate analysis of these samples \citep{Acacia2016}. We also use state-of-the-art models of the effect of non-linearities, bias and redshift-space distortions that allow us to extend our analysis of the full shape of the clustering wedges to smaller scales. We perform extensive tests of the performance of our methodology on $N$-body simulations and mock catalogues and find precise and accurate constraints. \begin{figure*} \includegraphics[width=0.98\textwidth]{figs/figure1.pdf} \caption{ Clustering wedges in the directions parallel (blue) intermediate (green) and transverse (red) to the line of sight measured from the combined galaxy sample of BOSS DR12 in our three redshift bins, as a function of the pair separation expressed in ${\rm Mpc}$ and $h^{-1}{\rm Mpc}$ in the lower and upper axes, respectively. The error bars correspond to the dispersion of the results inferred from a set of $N_{\rm m}=2045$ mock catalogues of the full BOSS survey. The solid lines correspond to the best-fitting model to these measurements obtained as described in Section~\ref{sec:bao}. } \label{fig:dr12_3w} \end{figure*} Our analysis is part of a series of papers examining the information in the anisotropic clustering pattern of the combined sample of BOSS DR12. \citet{Salazar2016} perform a tomographic analysis of the clustering properties of this sample by means of angular correlation functions in thin redshift shells. \citet{Grieb2016} use the same description of non-linearities, bias and RSD used in our analysis to extract cosmological information from the full shape of three clustering wedges measured in Fourier space. \citet{Satpathy2016} use a model based on convolution Lagrangian perturbation theory \citep{Carlson2013,Wang13} and the Gaussian streaming model \citep{Scoccimarro2004, Reid2011} to fit the full shape of the monopole and quadrupole of the two-point correlation function, $\xi_{0,2}(s)$. \citet{Beutler2016b} apply a model based on \citet{Taruya2010} to the power spectrum multipoles $P_{\ell}(k)$ for $\ell=0,2,4$. \citet{Tinker2016} present a comparison of the results of different RSD analysis techniques. \citet{Ross2016} and \citet{Beutler2016a} perform BAO-only fits to the Legengre multipoles of order $\ell=0,2$ of the two-point functions in configuration and Fourier space obtained after the application of the reconstruction technique \citep{Eisenstein2007b,Padmanabhan2012} as described in \citet{Cuesta2016}. The potential systematics of these BAO-only measurements are discussed in \citet{Vargas2016}. \citet{Acacia2016} use the methodology described in \citet{Sanchez2016} to combine the results presented here with those of the other full-shape and BAO-only analyses into a final set of BOSS consensus constraints and explore their cosmological implications. The outline of this paper is as follows. In Section~\ref{sec:boss} we describe our galaxy sample, the procedure we follow to measure the clustering wedges, and the mock catalogues used to compute our estimate of their covariance matrices. Our model of the full shape of the clustering wedges is described in Section~\ref{sec:model}, together with the tests we have performed by applying it to $N$-body simulations and mock catalogues. In Section~\ref{sec:results} we study the cosmological implications of our clustering measurements. After describing our methodology to obtain cosmological constraints in Section~\ref{sec:method}, Sections~\ref{sec:lcdm} to \ref{sec:gamma} describe the results we obtained from different combinations of data sets and parameter spaces. In Section~\ref{sec:bao} we compress the information of the BOSS clustering wedges into geometric constraints and measurements of the growth of structure. Finally, we present our main conclusions in Section~\ref{sec:conclusions}.
\label{sec:conclusions} We have analysed the cosmological implications of the measurements of three clustering wedges $\xi_{3{\rm w}}(s)$ of the final galaxy samples from BOSS corresponding to SDSS-DR12. We make use of the BOSS combined sample described in \citet{Reid2016}, containing the joint information of the LOWZ and CMASS samples that were analysed separately in former studies, including also the Early regions that were previously excluded. We have focussed on adjusting our analysis methodology to maximize the information extracted from the BOSS data. We implemented a state-of-the-art description of the effects of the non-linear evolution of density fluctuations, bias and RSD that allowed us to extract information from the full shape of our clustering measurements including smaller scales than in previous analyses. We performed extensive tests of this model using various N-body simulations and BOSS mock catalogues, showing that it can be used to extract cosmological information from our measurements of three clustering wedges for scales $s \gtrsim 20\,h^{-1}{\rm Mpc}$ without introducing any significant systematic errors. We used the information from our clustering measurements in combination with the latest CMB measurements from Planck and the JLA SN sample to constrain the parameters of the $\Lambda$CDM model and a number of its potential extensions, including more general dark energy models, non-flat universes, neutrino masses and possible deviations from the predictions of GR. Our results are completely consistent with the standard $\Lambda$CDM plus GR cosmological paradigm. When this model is extended by allowing one additional parameter to vary freely, the combination of the CMB data from Planck and our BOSS LSS measurements is enough to put tight constraints on the additional variable, with the SN data leading only to marginal improvements. The SN information is most useful when more than one additional parameter is included in the analysis, leading to final constraints in agreement with the canonical $\Lambda$CDM values. The full data set combination can constrain the dark energy equation of state parameter to $w_{\rm DE}=-0.996\pm0.042$ when assumed time-independent, with no indication of a departure from this value when it is allowed to evolve according to equation~(\ref{eq:wa}). The simultaneous variation of additional cosmological parameters does not affect this limit significantly. Our results are also completely consistent with the flat-Universe prediction from the most simple inflationary models, with $\Omega_{k}=-0.0007\pm 0.0030$. We derive tight constraints on the total sum of neutrino masses to $\sum m_{\nu} < 0.25\,{\rm eV}$ at 95 per cent CL. We also test the agreement of our clustering measurements with the predictions of GR by assuming the parametrization of equation~(\ref{eq:growth_gamma}) for the growth-rate of cosmic structure and find $\gamma = 0.609\pm 0.079$, in agreement with the GR value of $\gamma=0.55$. The information of our clustering measurements can be compressed into constraints on the parameter combinations $D_{\rm V}(z)/r_{\rm d}$, $F_{\rm AP}(z)$ and $f\sigma_8(z)$ at the mean redshifts of each of our three redshift bins with their respective covariance matrices. These results are in excellent agreement with the predictions of the best-fitting $\Lambda$CDM model to the CMB measurements from Planck, highlighting the consistency between these data sets. Our results are combined with those of our companion papers into a final set of consensus constraints in \citet{Acacia2016} using the methodology described in \citet{Sanchez2016}. Our results show that anisotropic clustering measurements have become one of the most powerful available cosmological probes. By exploiting the BAO and RSD signals imprinted in these measurements, the BOSS galaxy samples have significantly improved our knowledge of the basic cosmological parameters. The application of the methodology presented here to galaxy samples from future surveys such as the Dark Energy Spectroscopic Instrument \citep[DESI;][]{Levi:2013gra} and the ESA space mission \emph{Euclid} \citep{Laureijs:2011gra} will help to push our tests of the $\Lambda$CDM paradigm to even higher accuracies. A joint analysis of two-point statistics with higher-order measurements such as the three-point correlation function or the bispectrum \citep{GilMarin2015}, a detailed study of redshift-space distortions on small scales including the impact of effects such as velocity or assembly bias \citep{Reid2014}, or the advancement of methods to reconstruct the underlying density field \citep{Kitaura2016b} are strategies that could help to further increase the information extracted from LSS data sets, which will continue shaping our understanding of cosmic history
16
7
1607.03147
We explore the cosmological implications of anisotropic clustering measurements in configuration space of the final galaxy samples from Data Release 12 of the Sloan Digital Sky Survey III Baryon Oscillation Spectroscopic Survey. We implement a new detailed modelling of the effects of non-linearities, bias and redshift-space distortions that can be used to extract unbiased cosmological information from our measurements for scales s ≳ 20 h<SUP>-1</SUP> Mpc. We combined the information from Baryon Oscillation Spectroscopic Survey (BOSS) with the latest cosmic microwave background (CMB) observations and Type Ia supernovae samples and found no significant evidence for a deviation from the Λ cold dark matter (ΛCDM) cosmological model. In particular, these data sets can constrain the dark energy equation-of-state parameter to w<SUB>DE</SUB> = -0.996 ± 0.042 when to be assumed time independent, the curvature of the Universe to Ω<SUB>k</SUB> = -0.0007 ± 0.0030 and the sum of the neutrino masses to ∑m<SUB>ν</SUB> &lt; 0.25 eV at 95 per cent confidence levels. We explore the constraints on the growth rate of cosmic structures assuming f(z) = Ω<SUB>m</SUB>(z)<SUP>γ</SUP> and obtain γ = 0.609 ± 0.079, in good agreement with the predictions of general relativity of γ = 0.55. We compress the information of our clustering measurements into constraints on the parameter combinations D<SUB>V</SUB>(z)/r<SUB>d</SUB>, F<SUB>AP</SUB>(z) and fσ<SUB>8</SUB>(z) at z<SUB>eff</SUB> = 0.38, 0.51 and 0.61 with their respective covariance matrices and find good agreement with the predictions for these parameters obtained from the best-fitting ΛCDM model to the CMB data from the Planck satellite. This paper is part of a set that analyses the final galaxy clustering data set from BOSS. The measurements and likelihoods presented here are combined with others by Alam et al. to produce the final cosmological constraints from BOSS.
false
[ "SUB", "Baryon Oscillation Spectroscopic Survey", "cosmological model", "unbiased cosmological information", "the Sloan Digital Sky Survey III Baryon Oscillation Spectroscopic Survey", "BOSS", "good agreement", "cosmic structures", "CMB", "anisotropic clustering measurements", "configuration space", "SUB>k</SUB", "F<SUB", "general relativity", "Data Release", "Planck", "the latest cosmic microwave background (CMB) observations and Type Ia supernovae samples", "the final galaxy clustering data", "Alam et al", "the final cosmological constraints" ]
12.077656
2.499372
161
3893390
[ "Schnorr-Müller, Allan", "Davies, R. I.", "Korista, K. T.", "Burtscher, L.", "Rosario, D.", "Storchi-Bergmann, T.", "Contursi, A.", "Genzel, R.", "Graciá-Carpio, J.", "Hicks, E. K. S.", "Janssen, A.", "Koss, M.", "Lin, M. -Y.", "Lutz, D.", "Maciejewski, W.", "Müller-Sánchez, F.", "Orban de Xivry, G.", "Riffel, R.", "Riffel, Rogemar A.", "Schartmann, M.", "Sternberg, A.", "Sturm, E.", "Tacconi, L.", "Veilleux, S.", "Ulrich, O. A." ]
2016MNRAS.462.3570S
[ "Constraints on the broad-line region properties and extinction in local Seyferts" ]
48
[ "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany; CAPES Foundation, Ministry of Education of Brazil, 70040-020 Brasília, Brazil", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Department of Physics, Western Michigan University, Kalamazoo, MI 49008-5252, USA", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Department of Physics, Centre of Extragalactic Astronomy, Durham University, Durham DH1 3LE, UK", "Departamento de Astronomia, Universidade Federal do Rio Grande do Sul, IF, CP 15051, 91501-970 Porto Alegre, RS, Brazil", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Physics and Astronomy Department, University of Alaska, Anchorage, AK 99508-4664, USA", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Department of Physics, Institute for Astronomy, ETH Zurich, Wolfgang-Pauli-Strasse 27, CH-8093 Zurich, Switzerland", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Astrophysics Research Institute, Liverpool John Moores University, IC2 Liverpool Science Park, 146 Brownlow Hill, L3 5RF, UK", "Center for Astrophysics and Space Astronomy, University of Colorado, Boulder, CO 80309-0389, USA", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Departamento de Astronomia, Universidade Federal do Rio Grande do Sul, IF, CP 15051, 91501-970 Porto Alegre, RS, Brazil", "Departamento de Física, Centro de Ciências Naturais e Exatas, Universidade Federal de Santa Maria, 97105-900 Santa Maria, RS, Brazil", "Centre for Astrophysics and Supercomputing, Swinburne University of Technology, Hawthorn, Victoria 3122, Australia", "Raymond and Beverly Sackler School of Physics and Astronomy, Tel Aviv University, Ramat Aviv 69978, Israel", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstr. 1, D-85741 Garching, Germany", "Department of Astronomy and Joint Space-Science Institute, University of Maryland, College Park, MD 20742-2421, USA", "Department of Physics, Western Michigan University, Kalamazoo, MI 49008-5252, USA" ]
[ "2016ApJ...832....8B", "2017A&A...601A..32S", "2017A&A...607A..37M", "2017ApJ...836L...8C", "2017MNRAS.466.4917D", "2018A&A...615A.164V", "2018ARA&A..56..625H", "2018ApJ...856..154S", "2018MNRAS.473.2930D", "2018MNRAS.473.4582L", "2018RNAAS...2....3D", "2019A&A...628A..26P", "2019AJ....158..129K", "2019MNRAS.487.3884C", "2019arXiv191200164A", "2020A&A...634A.114C", "2020ApJ...888...71W", "2020ApJ...890L..29A", "2020ApJ...896..108Z", "2020MNRAS.498.4150D", "2020RAA....20..147X", "2021A&A...647A.195G", "2021A&A...648A.117G", "2021A&A...654A.132B", "2021ApJ...918...29M", "2021ApJ...922..179B", "2021IAUS..359..192L", "2021MNRAS.502.3618G", "2021MNRAS.508..144G", "2021RAA....21..199L", "2022A&A...666A...6W", "2022ApJ...929..184S", "2022ApJS..261....8R", "2022MNRAS.509..489L", "2022MNRAS.509.2377D", "2022MNRAS.516.2876M", "2023ApJ...950..106W", "2023ApJ...956...60C", "2023BAAA...64..247L", "2023MNRAS.518..130S", "2023MNRAS.520.1687A", "2023MNRAS.520.2781K", "2023Physi...5.1061D", "2024A&A...684A..52K", "2024MNRAS.tmp.1456M", "2024Univ...10...69D", "2024arXiv240414475B", "2024arXiv240419544B" ]
[ "astronomy" ]
7
[ "dust", "extinction", "galaxies: active", "galaxies: ISM", "galaxies: nuclei", "galaxies: Seyfert", "Astrophysics - Astrophysics of Galaxies" ]
[ "1970ApJ...160...25O", "1975MNRAS.171..395N", "1979RvMP...51..715D", "1980MNRAS.193..563W", "1981ApJ...249..462O", "1981ApJ...250..478K", "1983MNRAS.202..453B", "1985ApJ...289...67D", "1987ApJ...315...74W", "1988ApJ...332..172R", "1989ApJ...340..190G", "1992AJ....104.2072T", "1992ApJS...79...49W", "1992MNRAS.257..677W", "1993ApJ...404L..51N", "1994ApJ...428..124P", "1995A&A...293..889P", "1995ApJ...440..141G", "1995ApJ...448...27N", "1995ApJ...454...95M", "1995ApJ...455L.119B", "1996ApJS..104...37M", "1997ApJ...477..631V", "1997ApJS..108..401K", "1997MNRAS.291..403R", "1997MNRAS.292..273W", "1999A&A...346..764S", "2000A&A...355..485G", "2000ApJ...533..682C", "2000ApJ...536..284K", "2000ApJ...539..166R", "2000ApJ...539..718C", "2001A&A...379..125K", "2002A&A...384..826V", "2002ApJ...571..234R", "2002ApJ...572..746O", "2002ApJ...581..932B", "2004A&A...418..465L", "2004ApJ...606..749K", "2005ApJ...620..629D", "2005ApJ...630..122G", "2006ApJS..166..128Z", "2006agna.book.....O", "2007ApJ...671..104L", "2007arXiv0711.1013G", "2008A&A...490..995P", "2008AJ....135.2048T", "2008ApJ...685..160N", "2008MNRAS.383..581D", "2009ApJ...700.1173Z", "2009MNRAS.399.1293M", "2010ApJ...725.1749T", "2010ApJS..189...15K", "2010SPIE.7737E..28M", "2011A&A...525L...8C", "2011A&A...536A.105V", "2011MNRAS.410.1593W", "2012ApJ...745..107W", "2012ApJ...754...18R", "2012ApJ...759...95N", "2012MNRAS.420..526M", "2013A&A...558A.149B", "2013A&A...559A..96F", "2013RMxAA..49..137F", "2014ApJ...788...76W", "2014ApJ...792L...9L", "2014ApJ...795..164T", "2014MNRAS.438..604B", "2014MNRAS.439.1051L", "2014MNRAS.442.2145P", "2015A&A...578A..47B", "2015ARA&A..53..365N", "2015ApJ...806..127D", "2016A&A...586A..28B", "2016ApJ...832....8B", "2016MNRAS.461.4227H", "2017MNRAS.467..226G" ]
[ "10.1093/mnras/stw1865", "10.48550/arXiv.1607.07308" ]
1607
1607.07308_arXiv.txt
The broad emission lines observed in the spectra of type 1 AGNs, including the partially obscured intermediate type AGN in which at least some broad lines are seen \citep{Osterbrock70}, are the most direct tracers of the activity and immediate environment of accreting supermassive black holes. However, despite the prominence of these lines, their use in deriving the physical conditions of the broad line region (BLR) proved troublesome until the development of the locally optimally emitting cloud (LOC) model by \cite{baldwin95}. In this concept, the BLR comprises gas clouds covering a wide range of physical conditions; and line emission may arise from different regions of the ensemble according to where it is most optimally generated. The gas `clouds' in the LOC model do not necessarily correspond to actual clouds, but may simply be packets of gas within a cloud that itself comprises gas with variety of properties. As such, the more physically motivated models of the BLR, such as the radiation pressure confinement (RPC) model \citep{baskin14}, can also be considered LOC models. The most general set of LOC models has been explored and compared to reverberation mapping results by \cite{korista00} and \cite{korista04}; and applied to quasars by \cite{korista12}. Crucially, the generic models allow one to test the physical models such as RPC -- by indicating which part of the density vs photon-flux parameter space is consistent with the observed line emission, and comparing this to the specific prediction made by the physical model. Problems remain, however, in doing this because the intrinsic line ratios are not known {\it a priori} (the assumption of case B recombination as used for stellar photoionisation does not apply, primarily due to the high densities, high incident UV ionising fluxes, and optical depth of the emitted lines), and because of the intrinsic variability of the line emission on timescales of days to months. If the BLR is totally unobscured, then the observed line ratios are the same as the intrinsic line ratios. In this context, \cite{lamura07} selected 90 Seyfert\,1 galaxies from the Sloan Digital Sky Survey (SDSS) in which all broad lines from H$\alpha$ to H$\delta$ could be measured -- since this implied negligible extinction to the BLR. They derived $\langle$H$\alpha$/H$\beta\rangle = 3.45\pm0.65$, implying a wide range of intrinsic ratios, which they explored as a function of the line width. With a similar aim, \cite{dong08} selected 446 Seyfert\,1s from the SDSS which had blue optical continuum, which implies low extinction to the accretion disk. They found the distribution of H$\alpha$/H$\beta$ which could be well described by a log-Gaussian distribution peaking at 3.06 and with a standard deviation of only 0.03\,dex (corresponding to $\pm0.2$ in the ratio). This much tighter range of ratios led \cite{dong08} to suggest that one could adopt 3.06 as the intrinsic H$\alpha$/H$\beta$ and hence derive the extinction to the BLR. Recently, \citet{gaskell15} analysed a sample of low-redshift AGNs with extremely blue optical spectral indices and found that these objects have a mean ratio of 2.72\,$\pm$\,0.04. They interpreted this result as an indication that case B recombination is valid in the BLR, and that even the bluest 10\% of SDSS AGNs still have significant reddening. On the other hand, \citet{baron16} analysed 5000 SDSS type 1 AGNs at z$\approx$\,0.4 and while they found that the majority of the objects in their sample are reddened, in unreddened objects the mean H$\alpha$/H$\beta$ ratio is $\approx$\,3 with a broad distribution in the range 1.5--4, which suggests that simple Case B recombination cannot explain the ratio in all sources. Despite these results, significant uncertainties remain. These include the impact on the observed H$\alpha$/H$\beta$ ratio of low continuum fluxes in unreddened AGN \citep{osterbrock81,rudy88,trippe08} and also of even low extinction to the BLR \citep{goodrich89,tram92,goodrich95,dong05,trippe10}. Ideally one would like to derive the intrinsic line ratios and extinction simultaneously. \cite{korista12} took a step in this direction, using observations of optical and near-infrared broad lines in 4 optically selected quasars -- although these were not obtained simultaneously. These authors then compared the observed ratios to a grid of photoionisation models, in order to derive the best fitting density and photon flux. By focussing on ratios involving near-infrared lines (Pa$\alpha$, Pa$\beta$, Pa$\gamma$) as well as H$\alpha$, the impact of extinction could be reduced. In this paper we present high resolution spectra of Seyfert\,1s obtained simultaneously across the optical and infrared bands. For the first time we are able to analyse ratios involving up to 6 broad H\,I lines (H$\gamma$, H$\beta$, H$\alpha$, Pa$\beta$, Pa$\gamma$ and Pa$\delta$) in order to derive both the intrinsic line ratios and the extinction for individual objects, by comparison to a grid of photoionisation calculations. By using so many lines, we are also better able to account for the narrow line emission superimposed on the broad lines, in order to more robustly estimate the broad line fluxes. And by focussing here on the hydrogen recombination lines, for which the physical process is in principle relatively simple, we avoid complications due to, for example, metallicity \citep{korista12}. We note that additional problems associated with Ly$\alpha$ \citep{netzer95} are not relevant here since we do not make use of this line. In Section\,\ref{Observations} we describe the observations and data reduction. In Section\,\ref{sec:analysis} we describe the fitting of the broad and narrow emission lines, the photoionisation calculations and the extinction measurements. Then in Section\,\ref{sec:discuss} we discuss the results, and in Section\,\ref{sec:conc} we present our conclusions. \section {Observations and Data reduction}\label{Observations} \subsection{Sample Selection} The Seyfert 1 galaxies studied in this work are part of LLAMA (Local Luminous AGN with Matched Analogues), a study using VLT/SINFONI and VLT/Xshooter of a complete volume limited sample of nearby active galaxies selected by their 14--195\,keV luminosity. The rationale for the selection of the LLAMA sample is described in detail in \citet{davies15}. Briefly, these authors note that the accretion rate, traced by the AGN luminosity, is a crucial parameter that is often ignored, especially in AGN samples selected via line ratios. They argue that low luminosity AGN can in principle be powered by a very local gas reservoir on scales of $<10$\,pc; and it is only at higher luminosities that (observable) coherent dynamical mechanisms are required to drive sufficient gas inwards from circumnuclear scales of 100-1000\,pc. If one wishes to study such mechanisms, the target galaxies must be near enough that these scales can be very well spatially resolved. While there are many tracers that can be used to select local AGN based on luminosity, the very hard X-rays ($>10$\,keV) are widely accepted to be among the least biassed with respect to host galaxy properties. The all-sky {\it Swift BAT} survey at 14-195\,keV therefore provides an excellent parent sample for the selection. By selecting appropriate luminosity and distance (redshift) thresholds ($\log{L_{14-195keV} [erg s^{-1}]} > 42.5$ and $z < 0.01$) one can select a complete volume limited sample from the flux limited survey. An additional requirement that the objects are observable from the VLT lead to a sample of 20 AGN, of which 10 are classified as type 1.0-1.9 \citep{davies15}. \subsection{Observations} \begin{table*} \begin{center} \begin{tabular}{c c c c c c c c c} \hline \hline Object & Date(s) & Exposure & Air Mass & Seeing & AGN & log\,L$_{14-195}$ &log\,N$_{H}$ &Distance\\ & Observed & (s) & & (\arcsec) & Classification & (erg\,s$^{-1}$) & (cm$^{-2}$) & (Mpc)\\ \hline \multirow{2}{*}{MCG0514} & 23 Nov 2013 & 1920 & 1.012 & 0.61 & \multirow{2}{*}{Sy 1.0} & \multirow{2}{*}{42.60} & \multirow{2}{*}{$\le21.9^{1}$} & \multirow{2}{*}{41}\\ & 11 Dec 2013 & 1920 & 1.002 & 0.72 & & & & \\ MCG0523 & 22 Jan 2014 & 3840 & 1.063 & 1.21 & Sy 1.9 & 43.47 & 22.1-22.7$^2$ & 35\\ MCG0630 & 16 Jun 2015 & 1920 & 1.124 & 0.83 & Sy 1.2 & 42.74 & 20.9$^1$-22.2$^3$ & 27\\ NGC1365 & 10 Dec 2013 & 1920 & 1.022 & 1.34 & Sy 1.8 & 42.39 & 22.0-23.4$^{4}$ & 18\\ \multirow{2}{*}{NGC2992} & 26 Feb 2014 & 3840 & 1.302 & 0.72 & \multirow{2}{*}{Sy 1.8} & \multirow{2}{*}{42.62} & \multirow{2}{*}{21.6-22.2$^2$} & \multirow{2}{*}{36}\\ & 27 Feb 2014 & 3840 & 1.342 & 0.72 & & & &\\ NGC3783 & 11 Mar 2014 & 3840 & 1.428 & 0.81 & Sy 1.2 & 43.49 & 20.5$^1$-22.5$^5$ & 38\\ NGC4235 & 13 May 2015 & 1920 & 1.203 & 0.73 & Sy 1.2 & 42.72 & 21.2$^6$-21.3$^1$ & 37\\ NGC4593 & 10 Mar 2014 & 1920 & 1.325 & 0.80 & Sy 1.2 & 43.16 & 20.1$^7$-20.3$^5$ & 37\\ NGC6814 & 13 May 2015 & 3840 & 1.077 & 0.86 & Sy 1.2 & 42.69 & 21.0$^1$-21.1$^3$ & 23\\ \hline \end{tabular} \caption{Summary of the observations, Seyfert sub-types (see section\,3.4), 14-195\,kev luminosity (taken from \citealt{davies15}), X-ray absorbing column (for objects which show N$_{H}$ variations we show the smallest and largest values in the literature) and distance (taken from NED, using redshift independent estimates). \textbf{Abbreviated names}:MCG--05-14-012, MCG--05-23-16 and MCG--06-30-015. \textbf{References:}$^{1}$:\citet{davies15},$^{2}$:\citet{risaliti02},$^{3}$:\citet{molina09},$^{4}$:\citet{walton14},$^{5}$:\citet{lutz04},$^{6}$:\citet{papadakis08},$^{7}$:\citet{winter12}.} \label{table_obj} \end{center} \end{table*} High spectral resolutions were obtained using the X-shooter spectrograph mounted at the ESO-VLT in Paranal, Chile, between the nights of November 23, 2013 and June 16, 2015. A summary of the observations and basic properties of the sources are presented in Table~\ref{table_obj}. X-shooter is a multi wavelength high resolution spectrograph which consists of 3 spectroscopic arms, covering the 3000--5595\r{A} (UVB arm), 5595--10240\r{A} (VIS arm) and 10240--24800\r{A} (NIR arm) wavelength ranges. Each arm consists of an independent cross dispersed echelle spectrograph with its optimised optics, dispersive element and detector \citep{vernet11}. The observations were performed in IFU-offset mode, with a field-of-view of 1\farcs8$\,\times$\,4\arcsec. In IFU mode the resolving power of each arm is R\,$\approx$\,8400 for the UVB arm, R\,$\approx$\,13200 for the VIS arm and R\,$\approx$\,8300 for the NIR arm. The observations of each target consisted of two observing blocks with sky observations performed after each target observation. A spectroscopic standard star was observed on the same night as the target observation under similar atmospheric conditions. Telluric standard star observations were performed both before and after the target observations. \subsection{Data Reduction} The X-shooter spectra were reduced within the ESO Reflex environment \citep{freudling13}, using version 2.6.0 of the ESO X-shooter pipeline in IFU/Offset mode \citep{modigliani10}, in order to produce a 1\farcs8$\,\times$\,4\arcsec\ datacube. The reduction routine comprised detector bias and dark current subtraction, and subsequent rectification of the spectra and wavelength calibration. Relative and absolute flux calibration of the datacube were performed using data of a spectroscopic standard star observed on the same night as the source. We corrected the spectra for telluric absorption using a telluric standard star at a similar air mass and observed either before or right after the target observation. A more detailed discussion of the specific issues encountered when processing these data, and of their solutions, will be given in Burtscher et al. (in prep).
\label{sec:conc} In this paper, we present simultaneous spectroscopic observations of 6 optical (Balmer) and near-infrared (Paschen) broad lines from nine nearby Seyferts that are part of our ongoing survey of local luminous AGN with matched analogs (LLAMA). Analysing the H\,I line ratios in the context of a grid of photoionisation models (and without assuming case B recombination), we are able to derive the intrinsic line ratios, photon flux and density of the BLR, as well as the extinction to it. Based on the analysis of these objects, our main conclusions are: \begin{itemize} \item Values of the intrinsic H$\alpha$/H$\beta$ line ratio lie in the range 2.5--6.6 (without accounting for uncertainties), confirming that case B recombination cannot be assumed, and that a single set of excitation properties cannot be used to represent all BLRs. \item The density of the H\,I emitting gas in the BLR is typically 10$^{11}$\,cm$^{-3}$ for the sample, although there may be variation between individual sources. The ionising photon flux is in the range $\sim10^{17.5}$--10$^{18.5}$\,cm$^{-2}$\,s$^{-1}$, which implies that a significant amount of the H\,I emitting gas is located at regions near the dust sublimation radius. These values are consistent with theoretical predictions, in particular with the radiation pressure confinement model. \item The Seyfert sub-types, determined via the [O\,III]/H$\beta$ line ratio, are consistent with the extinction to the BLR, which is based on independent estimates from the H\,I and He\,II lines: in Seyfert 1.0 and 1.2s the BLR is either unobscured or mildly obscured while in Seyfert 1.8 and 1.9s the BLR is moderately obscured with extinction in the range A$_V = 4$--8\,mag. By inference, Seyfert~2s have A$_V > 8$\,mag. \item In the moderately obscured objects the extinction to the BLR is significantly larger than the extinction to the NLR. This may be caused by the torus but could also be due to dusty filaments or a nuclear spiral. In the specific case of MCG--05-23-16, the nuclear filament cannot account for the obscuration, which is instead likely due to the torus. \end{itemize}
16
7
1607.07308
We use high-spectral resolution (R &gt; 8000) data covering 3800-13 000 Å to study the physical conditions of the broad-line region (BLR) of nine nearby Seyfert 1 galaxies. Up to six broad H I lines are present in each spectrum. A comparison - for the first time using simultaneous optical to near-infrared observations - to photoionization calculations with our devised simple scheme yields the extinction to the BLR at the same time as determining the density and photon flux, and hence distance from the nucleus, of the emitting gas. This points to a typical density for the H I emitting gas of 10<SUP>11</SUP> cm<SUP>-3</SUP> and shows that a significant amount of this gas lies at regions near the dust sublimation radius, consistent with theoretical predictions. We also confirm that in many objects, the line ratios are far from case B, the best-fitting intrinsic broad-line Hα/H β ratios being in the range 2.5-6.6 as derived with our photoionization modelling scheme. The extinction to the BLR, based on independent estimates from H I and He II lines, is A<SUB>V</SUB> ≤ 3 for Seyfert 1-1.5s, while Seyfert 1.8-1.9s have A<SUB>V</SUB> in the range 4-8. A comparison of the extinction towards the BLR and narrow-line region (NLR) indicates that the structure obscuring the BLR exists on scales smaller than the NLR. This could be the dusty torus, but dusty nuclear spirals or filaments could also be responsible. The ratios between the X-ray absorbing column N<SUB>H</SUB> and the extinction to the BLR are consistent with the Galactic gas-to-dust ratio if N<SUB>H</SUB> variations are considered.
false
[ "BLR", "gas", "NLR", "scheme", "Seyfert", "theoretical predictions", "regions", "photoionization calculations", "dust", "N", "independent estimates", "scales", "the line ratios", "our devised simple scheme", "the emitting gas", "N<SUB", "case B", "SUP>-3</SUP", "the BLR and narrow-line region", "the dust sublimation radius" ]
15.406553
7.690591
119
12593903
[ "Spinoso, Daniele", "Bonoli, Silvia", "Dotti, Massimo", "Mayer, Lucio", "Madau, Piero", "Bellovary, Jillian" ]
2017MNRAS.465.3729S
[ "Bar-driven evolution and quenching of spiral galaxies in cosmological simulations" ]
74
[ "Università degli Studi di Milano-Bicocca, Piazza della Scienza 3, I-20126 Milano, Italy; Centro de estudios de fisí ca del cosmos de Aragó n, plaza San Juan, 1 planta-2 E-44001 Teruel, Spain", "Centro de estudios de fisí ca del cosmos de Aragó n, plaza San Juan, 1 planta-2 E-44001 Teruel, Spain", "Università degli Studi di Milano-Bicocca, Piazza della Scienza 3, I-20126 Milano, Italy; INFN, Sezione di Milano-Bicocca, Piazza della Scienza 3, I-20126 Milano, Italy", "Center for Theoretical Astrophysics and Cosmology, Institute for Computational Science, University of Zurich, Winterthurerstrasse 190, CH-8057 Z¨urich, Switzerland; Kavli Institute for Theoretical Physics, Kohn Hall, University of California, Santa Barbara, CA 93106-4030, USA", "Department of Astronomy and Astrophysics, University of California, 1156 High Street, Santa Cruz, CA 95064, USA", "Department of Astrophysics, American Museum of Natural History, Central Park West and 79th St, New York, NY 10024, USA" ]
[ "2017ApJ...835..289S", "2017ApJ...838..105L", "2017MNRAS.469.2806A", "2017MNRAS.469.3722R", "2017PASA...34...41N", "2018A&A...609A..60K", "2018ApJ...868..100G", "2018MNRAS.473.2608Z", "2018MNRAS.473.4731K", "2018MNRAS.474.1909F", "2018MNRAS.474.3101J", "2018MNRAS.476.2318S", "2018MNRAS.478.1231P", "2018MNRAS.479.5214Z", "2018MNRAS.480.1340I", "2019A&A...621L...4G", "2019A&A...628A..24G", "2019ApJ...870...19G", "2019ApJ...874...67B", "2019MNRAS.485.5073D", "2019MNRAS.487.3102G", "2019MNRAS.488..609I", "2019MNRAS.488.1864Z", "2019MNRAS.488L...6F", "2019MNRAS.489.3553E", "2019PASJ...71S..14S", "2020A&A...644A..79G", "2020ApJ...895...92Z", "2020MNRAS.491.1800B", "2020MNRAS.491.2547R", "2020MNRAS.491.3266G", "2020MNRAS.492.4697N", "2020MNRAS.494.5839K", "2020MNRAS.495.4681I", "2020MNRAS.499.1116F", "2020RAA....20...15A", "2021A&A...651A.107G", "2021A&A...654A.135D", "2021A&A...656A.133Q", "2021ApJ...911...57Z", "2021ApJ...913...37C", "2021ApJ...916...38Z", "2021MNRAS.500..282L", "2021MNRAS.502.2238M", "2021MNRAS.503.4992F", "2021MNRAS.507.4389G", "2022A&A...668L...3L", "2022ApJ...926...81A", "2022ApJ...940....1T", "2022MNRAS.509..567S", "2022MNRAS.510.5164C", "2022MNRAS.512.5339R", "2022MNRAS.513.2850B", "2022MNRAS.513.3768I", "2022MNRAS.514.1006I", "2022MNRAS.515..271R", "2023A&A...673A.147P", "2023ApJ...944L..25H", "2023ApJ...945L..10G", "2023ApJ...949...91Z", "2023MNRAS.518...13H", "2023MNRAS.519.4801C", "2023MNRAS.521.1775G", "2023MNRAS.523.5823B", "2024A&A...684A.179R", "2024ApJ...964..120P", "2024ApJ...968...87K", "2024FPP....1100059F", "2024MNRAS.528.6768H", "2024MNRAS.529..979L", "2024MNRAS.530.1171C", "2024MNRAS.tmp.1510S", "2024arXiv240411656S", "2024arXiv240505960G" ]
[ "astronomy" ]
15
[ "methods: numerical", "galaxies: bulges", "galaxies: evolution", "galaxies: formation", "galaxies: kinematics and dynamics", "galaxies: structure", "Astrophysics - Astrophysics of Galaxies" ]
[ "1964ApJ...139.1217T", "1976ApJ...209...53S", "1977ApJ...217..916S", "1978ApJ...222..850G", "1979ApJ...233...67R", "1979ApJ...233..857G", "1981A&A....96..164C", "1989Natur.338...45S", "1992MNRAS.259..345A", "1993A&A...271..391C", "1993RPPh...56..173S", "1997A&A...323..363M", "1997ApJ...482L.135M", "1997ApJ...487..591H", "1998ApJ...499..149M", "1998MNRAS.300...49B", "1999ApJ...510..125S", "1999ApJ...516..660H", "1999MNRAS.304..475M", "2000A&A...362..435L", "2000ApJ...529...93K", "2001sac..conf..223C", "2002ApJ...567...97L", "2002MNRAS.330...35A", "2004AJ....127..279B", "2004ApJ...600..595R", "2004ApJ...607..103L", "2004NewA....9..137W", "2005AJ....130..506B", "2005ApJ...630..837J", "2005MNRAS.363..496A", "2006ApJ...637..214M", "2006ApJ...645..209D", "2007ApJ...666..189B", "2007ApJS..170..377S", "2007PASA...24..159P", "2008ApJ...687L..13R", "2008MNRAS.388.1803R", "2008gady.book.....B", "2009ApJ...697..293D", "2010ApJ...714L.260N", "2010ApJ...719.1470V", "2010ApJ...721L.148B", "2010MNRAS.406.2267F", "2010MNRAS.408..812S", "2011ApJ...733L..47F", "2011ApJ...742...76G", "2011MSAIS..18..185G", "2012ASPC..453..289M", "2012ApJ...745..125L", "2012ApJ...750..141L", "2012ApJ...754L..29W", "2012ApJ...757...60K", "2012ApJ...758...14K", "2012ApJ...760...50S", "2012MNRAS.425L..10S", "2013A&A...549A.141A", "2013ApJ...765...89S", "2013ApJ...772...36G", "2013ApJ...773...43B", "2013ApJ...773L..32R", "2013ApJ...776...50C", "2013ApJ...779..162C", "2013MNRAS.429.1949A", "2013seg..book....1K", "2014A&A...561A..86M", "2014ApJ...795..104W", "2014MNRAS.441.1615G", "2014MNRAS.443.1125T", "2014MNRAS.445.3352C", "2014RvMP...86....1S", "2015A&A...575A..74S", "2015A&A...576A..16G", "2015A&A...576A.102A", "2015A&A...579A...2I", "2015A&A...580A.116G", "2015ApJ...800...19R", "2015ApJ...801...80L", "2015MNRAS.446.1957F", "2015MNRAS.446.2468E", "2015MNRAS.447..506C", "2015MNRAS.447.1774G", "2015MNRAS.450.2514I", "2015MNRAS.454.3641F", "2015PASJ...67...63O", "2016A&A...588A..33Q", "2016A&A...589A..35S", "2016A&A...591A..38C", "2016ApJ...819....1I", "2016ApJ...826..227L", "2016MNRAS.456.2848H", "2016MNRAS.459.2603B", "2016MNRAS.462.3727P" ]
[ "10.1093/mnras/stw2934", "10.48550/arXiv.1607.02141" ]
1607
1607.02141_arXiv.txt
Bars are extremely common non axisymmetric features in disk galaxies, occurring in up to $\gsim 30\%$ of massive ($M_* \gsim 10^{9.5} \msun$) spirals in the local Universe \citep{Laurikainen04,Nair10,Lee12a,gavazzi15}. Bars are considered to play a key role in the evolution of disk galaxies, being able to drive strong inflows of gas towards the central galactic regions \citep[e.g.][]{sand76, Roberts79, Athanassoula92} and triggering nuclear star-bursts \citep[e.g.][]{Ho97, Martinet97, Hunt99, Laurikainen04, Jogee05}. Bars are also thought to be responsible for the build-up of the pseudo/disky bulges, whose nearly exponential profiles hints to a disk origin \citep[e.g.][for a review]{kor13}. These structures are the most common type of bulges among galaxies in the stellar-mass range $10^{9.5}\msun <M_*<10^{10.5} \msun$, while classical bulges dominate among more massive systems \citep[e.g.][]{FisherDrory11}. Bars can also be responsible for triggering AGN activity, if a fraction of the inflowing gas can reach the central sub-pc of the galaxy \citep[e.g.][]{Shlosman89, Berentzen98, sellwood99, comb00, quere15, fanali15}. On longer timescales, the removal of the gas forced towards nuclear regions affects the star formation processes within the bar extent \citep{cheung13, fanali15}, contributing to the lowering of the specific star formation rate in the most massive spiral galaxies at low redshift \citep{cheung13, gavazzi15}. In addition to the effect of the bar onto the inter stellar medium (ISM), the dynamical evolution of the bar itself is advocated to be responsible for the boxy/peanut-shaped stellar bulges (B/P bulges hereafter) \citep[see][for a review]{kor13, Sellwood14}, observed in $\gsim 40 \%$ of edge on disk galaxies \citep[e.g.][]{Lutticke00}. Together with the high fraction of disky pseudobulges, the frequent occurrence of B/P bulges hints at the fundamental importance of secular evolution in the shaping of the central regions of disk galaxies. Most of the theoretical studies that support the existence of causal connections between bars and the above-mentioned structures/processes are either analytical or based on simulations of isolated galaxies \citep[e.g.][]{Athanassoula92,Berentzen98,Regan04,Berentzen07,VillaVargas10,Kim12,Cole14}. Although these kind of simulations are extremely informing about the dynamical effect of bars, cosmological simulations are needed to follow bar formation within the hierarchical growth of galaxies \citep[as discussed in, e.g.,][]{kor13}. Furthermore, most of the simulation literature on bar formation and evolution is based on collisionless simulations. A few works have employed hydrodynamics and star formation, showing interesting differences on important issues such as bar survival and bar-buckling \citep[see e.g.][]{deb06, athan13, athan05}, but none of them has employed modern sub-grid recipes for feedback, which constitute a crucial aspect of recent progress in simulating galaxy formation. To date only an handful of fully cosmological simulations have achieved the required numerical resolution and included all the physical processes needed to self-consistently produce barred galaxies \citep[e.g.][]{RomanoDiaz08,Scannapieco12,Kraljic12,Goz14,bonol15, fiac15, okamoto15}. Among the above-mentioned cosmological simulations of barred disk galaxies, \erisbh \citep{bonol15} and Argo \citep{fiac15} share the highest spatial and mass resolutions\footnote{The simulations by \cite{Kraljic12} have a comparable spatial resolution but a coarser resolution in mass.}, but Argo has been evolved only down to $z=3$, while \erisbh has been followed down to $z=0$, so its properties can be compared directly with the observed properties of well-resolved barred galaxies. \erisbh is a twin simulation of Eris \citep{gued11}, with which it shares initial conditions, resolution and sub-grid physics, but, unlike Eris, it also includes prescriptions for the formation, growth and feedback of supermassive black holes. Both Eris and \erisbh reseamble, at $z=0$, a late-type galaxy such as the Milky Way \citep{gued11, bonol15}, but while Eris hosts a typical pseudobulge, \erisbh features a strong bar and its bulge has a clear boxy-peanut morphology \citep{bonol15}. The aim of this work is to study the buildup and the evolution of the strong bar seen in \erisbh, to learn about the origin of bars and the impact that these structures have in shaping galaxies like our own Milky Way. The paper is organized as follows. In Section \ref{simulation} we briefly summarize the properties and main results of the \erisbh simulation. In Section \ref{sfc_dnst} we study the buildup of the bar, quantifying its strength and radial extent; we analyze the dynamical properties of the galaxy disk, testing its stability to non-axisymmetric perturbations and looking for resonances between the bar bulk precession and the orbital motions of disk stars; we also analyze the formation of the B/P morphology of the bulge. In Section \ref{gas_response} we show the impact of the bar in depleting gas in and triggering star formation in the central region of the galaxy. Finally, in Section \ref{conclusions} we summarize and discuss our results.
\label{conclusions} We analyzed the high-resolution cosmological \erisbh run \citep{bonol15}, which closely resembles an Sb/Sc galaxy with stellar mass and rotation velocity comparable to our Milky Way. At $z=0$ the galaxy forming at the centre of the refined region features a strong nuclear ($R\approx 2$ kpc) bar which is able to strongly influence: $(i)$ the dynamics of the stellar disk, including the formation of a B/P bulge in its centre; $(ii)$ the dynamics of the gas within the central $3$ kpc, which falls towards the galactic centre triggering a short burst of star formation in the galactic nucleus (within $\sim 600$ pc) as soon as the bar starts growing; $(iii)$ the late star formation in the central $\sim 3$ kpc. This is the consequence of the fast gas removal operated by the bar preventing any strong star formation episode after its formation. The analysis of the torques operated by the bar supports the notion that the bar efficiently drives gas inflows down to the resolution limit ($\sim$ hundreds of pc, due to the absence of any clear ILR) at any $z\lsim 0.4$. The absence of an intense star formation activity in the central regions of the disk as well as of strong AGN activity is purely due to the absence of dense gas within the bar extent due to rapid consumption by star formation at the onset of bar formation. The lack of a clear observational correlation between AGN activity and the occurrence of bars in galaxies \citep[see e.g.][for the different point of views]{Ho97, Mulchaey97, Hunt99, Knapen00, Laine02, Lee12b, Alonso13, Cisternas13, Cheung15} could be related to the prompt removal of gas. If we assume the results of \erisbh appy to the whole class of field disk galaxies in low density environments, we argue that the stronger gas inflow and enahnced star formation happen at the onset of bar formation, when the detection of a bar is more difficult as the bar is shorter and less regular in shape. Instead, when the bar is stronger and well-developed, hence easily determined from photometry or imaging, star formation has already ceased creating a ``dead zone'' in the galactic centre and making the occurrence of any nuclear activity less probable \citep[see e.g. the discussion in][]{fanali15}. Strong bars may arise at earlier times in more massive galaxies or galaxies living in dense environments, which evolve on shorter dynamical timescales. Hence we argue that bar formation can contribute to quenching and the formation of ``red nuggets'' at $z > 1$, as also suggested by the results of the ARGO simulations which exhibit several example of early bar formation leading to increased central baryonic densities (Fiacconi, Feldmann \& Mayer 2015). Bar-driven quenching should thus be seen as an alternative to mergers, disk fragmentation into massive clumps and AGN feedback, the main mechanisms explored in the literature over the last few years. Of course bar-driven quenching is related to feedback mechanisms operating in the central region, as it seems to be the case in \erisbh where AGN feedback might be instrumental in creating favourable conditions for bar formation at later stages. Since bar-formation requires a kinematically cold, thin disk to occur, it remains to be seen if this can be achieved by the latest generation of strong feedback models adopted in galaxy formation simulations. It is interesting to note that such a strong bar is absent in the Eris run, which differs from \erisbh only because it does not feature any MBH accretion and feedback prescription. This would seem to be at odd with the limited gas accretion occurring onto the central MBH \citep{bonol15}, that would imply a moderate effect of AGN feedback onto the host galaxy. However, at $z > 1$ there are transient near-Eddington accretion phases which ought to have an effect on the build-up of the central baryonic distribution. Indeed at $z < 1$ \erisbh has a much flatter rotation curve near the center as a result of the suppressed growth of the central baryonic density. The actual trigger of bar growth is still to be pinpointed. The main galaxy in the \erisbh run becomes bar unstable at large redshift (see Figure~\ref{growth}), but the bar structure forms only after the last minor merger episode. As discussed in section~\ref{sfc_dnst}, the properties of the bar do resemble those predicted for a tidally induced one. Whether the merger itself does provide the trigger for the instability to grow is unclear, as it is impossible to definitively constrain the time in between the merger and the actual onset of the bar growth. In order to test the possible tidal nature of the bar we plan to run a set of simulations restarting the \erisbh run before the merger, removing the particles forming the satellite, and checking whether the bar grows regardless of the perturbation. In conclusion, the present analysis of the \erisbh run has demonstrated that a bar resulting from the fully cosmological evolution of a% disk galaxy with quiet merger history strongly affects its host, in particular by quenching its star formation on kpc scales. This result provides further theoretical support to the recent claim by \cite{gavazzi15} that bars are one of the main contributor of the flattening observed at high masses in the star formation rate-stellar mass correlation \citep{whitaker12, magnelli14, whitaker14, gavazzi15b, ilbert15, lee15, schreiber16}.
16
7
1607.02141
We analyse the outputs of the cosmological 'zoom-in' hydrodynamical simulation ErisBH to study a strong stellar bar which naturally emerges in the late evolution of the simulated Milky Way-type galaxy. We focus on the analysis of the formation and evolution of the bar and on its effects on the galactic structure, the gas distribution and the star formation. A large central region in the ErisBH disc becomes bar unstable after z ∼ 1.4, but a clear bar starts to grow significantly only after z ≃ 0.4, possibly triggered by the interaction with a massive satellite. At z ≃ 0.1, the bar stabilizes and reaches its maximum radial extent of l ≈ 2.2 kpc. As the bar grows, it becomes prone to buckling instability. The actual buckling event, observable at z ≃ 0.1, results in the formation of a boxy-peanut bulge clearly discernible at z = 0. During its early growth, the bar exerts a strong torque on the gas and drives gas inflows that enhance the nuclear star formation on sub-kpc scales. Later on, as the bar reaches its maximum length and strength, the gas within its extent is nearly all consumed into stars, leaving behind a gas-depleted region in the central ∼2 kpc. Observations would more likely identify a prominent, large-scale bar at the stage when the galactic central region has already been gas depleted, giving a hint at the fact that bar-driven quenching may play an important role in the evolution of disc-dominated galaxies.
false
[ "bar", "gas inflows", "sub-kpc scales", "stars", "z", "evolution", "buckling instability", "bar-driven quenching", "a strong stellar bar", "Milky Way", "the nuclear star formation", "disc-dominated galaxies", "the galactic central region", "a clear bar", "the central ∼2 kpc", "the star formation", "l ≈ 2.2 kpc", "the simulated Milky Way-type galaxy", "a gas-depleted region", "A large central region" ]
10.425327
6.72554
133
3481761
[ "Lange, Rebecca", "Moffett, Amanda J.", "Driver, Simon P.", "Robotham, Aaron S. G.", "Lagos, Claudia del P.", "Kelvin, Lee S.", "Conselice, Christopher", "Margalef-Bentabol, Berta", "Alpaslan, Mehmet", "Baldry, Ivan", "Bland-Hawthorn, Joss", "Bremer, Malcolm", "Brough, Sarah", "Cluver, Michelle", "Colless, Matthew", "Davies, Luke J. M.", "Häußler, Boris", "Holwerda, Benne W.", "Hopkins, Andrew M.", "Kafle, Prajwal R.", "Kennedy, Rebecca", "Liske, Jochen", "Phillipps, Steven", "Popescu, Cristina C.", "Taylor, Edward N.", "Tuffs, Richard", "van Kampen, Eelco", "Wright, Angus H." ]
2016MNRAS.462.1470L
[ "Galaxy And Mass Assembly (GAMA): M_star - R_e relations of z = 0 bulges, discs and spheroids" ]
102
[ "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia", "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia", "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia; Scottish Universities' Physics Alliance (SUPA), School of Physics and Astronomy, University of St. Andrews, North Haugh, St. Andrews KY16 9SS, UK", "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia", "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia; Australian Research Council Centre of Excellence for All-sky Astrophysics (CAASTRO), 44 Rosehill Street, Redfern, NSW2016, Australia", "Astrophysics Research Institute, Liverpool John Moores University, IC2, Liverpool Science Park, 146 Brownlow Hill, Liverpool L3 5RF, UK", "School of Physics and Astronomy, The University of Nottingham, University Park, Nottingham NG7 2RD, UK", "School of Physics and Astronomy, The University of Nottingham, University Park, Nottingham NG7 2RD, UK", "NASA Ames Research Center, N232, Moffett Field, Mountain View, CA 94035, USA", "Astrophysics Research Institute, Liverpool John Moores University, IC2, Liverpool Science Park, 146 Brownlow Hill, Liverpool L3 5RF, UK", "Sydney Institute for Astronomy, School of Physics A28, University of Sydney, NSW 2006, Australia", "Astrophysics Group, School of Physics, University of Bristol, Bristol BS8 1TL, UK", "Australian Astronomical Observatory, PO Box 915, North Ryde, NSW 1670, Australia", "Department of Physics and Astronomy, University of the Western Cape, Robert Sobukwe Road, Bellville 7535, South Africa", "Research School of Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611, Australia", "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia", "European Southern Observatory, Alonso de Cordova 3107, Vitacura, Santiago, Chile", "University of Leiden, Sterrenwacht Leiden, Niels Bohrweg 2, NL-2333 CA Leiden, the Netherlands", "Sydney Institute for Astronomy, School of Physics A28, University of Sydney, NSW 2006, Australia", "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia", "School of Physics and Astronomy, The University of Nottingham, University Park, Nottingham NG7 2RD, UK", "Universität Hamburg, Hamburger Sternwarte, Gojenbergsweg 112, D-21029 Hamburg, Germany", "Astrophysics Group, School of Physics, University of Bristol, Bristol BS8 1TL, UK", "Jeremiah Horrocks Institute, University of Central Lancashire, Leighton Building, Preston PR1 2HE, UK; The Astronomical Institute of the Romanian Academy, Str. Cutitul de Argint 5, Bucharest, Romania", "School of Physics, the University of Melbourne, VIC 3010, Australia", "Max-Planck-Institut für Kernphysik, Saupfercheckweg 1, D-69117 Heidelberg, Germany", "European Southern Observatory, Karl-Schwarzschild-Str.2, D-85748 Garching, Germany", "International Centre for Radio Astronomy Research, University of Western Australia, M468, 35 Stirling Highway, Crawley, WA 6009, Australia" ]
[ "2016MNRAS.462.4336M", "2017AJ....153..111G", "2017ApJ...849...55S", "2017MNRAS.466.1513R", "2017MNRAS.466.3569P", "2017MNRAS.467..179G", "2017MNRAS.467.1033B", "2017MNRAS.467.2066S", "2017MNRAS.467.3751B", "2017MNRAS.469..968X", "2017MNRAS.471..447S", "2017PASA...34...47D", "2018ApJ...863...21Y", "2018ApJS..234...21D", "2018MNRAS.474..898A", "2018MNRAS.475..894T", "2018MNRAS.477.4711G", "2018MNRAS.478.4952F", "2018MNRAS.478.5410D", "2018MNRAS.479.3076A", "2018MNRAS.481.1376Z", "2018MNRAS.481.3573L", "2018NatAs...2..483V", "2019A&A...626A.110B", "2019A&A...629A..59P", "2019A&A...630A.113B", "2019A&A...632A..15L", "2019ApJ...886L..28M", "2019MNRAS.483.1881D", "2019MNRAS.483.2424R", "2019MNRAS.484..869V", "2019MNRAS.485.1477K", "2019MNRAS.487.1808M", "2019MNRAS.487.5416T", "2019MNRAS.488.1941M", "2019MNRAS.489.2830M", "2019MNRAS.489.4135D", "2019MNRAS.490.4060C", "2019RAA....19....6Z", "2020A&A...634A..11D", "2020A&A...641A.119B", "2020ApJ...897..102C", "2020ApJ...900..178L", "2020ApJ...905..167S", "2020IAUS..353..213V", "2020MNRAS.495.2827C", "2020MNRAS.496.3235B", "2020MNRAS.497.2786T", "2020MNRAS.499.1948L", "2020MNRAS.499.3399R", "2021ApJ...911...21B", "2021ApJ...923..205Y", "2021MNRAS.501.4359Z", "2021MNRAS.502.2446E", "2021MNRAS.502.3101N", "2021MNRAS.502.5370W", "2021MNRAS.504.3058M", "2021MNRAS.505.2247C", "2021MNRAS.505.3078V", "2021MNRAS.506.4011E", "2021Sci...372.1201T", "2021gamo.book.....H", "2022A&A...660A..20Z", "2022A&A...664A..92H", "2022ApJ...925..183H", "2022ApJ...929...61D", "2022ApJ...929..152L", "2022MNRAS.509.4372L", "2022MNRAS.510.3967R", "2022MNRAS.511.5475M", "2022MNRAS.513..256D", "2022MNRAS.513.2985R", "2022MNRAS.513.3709C", "2022MNRAS.514.6120J", "2022MNRAS.515.1175H", "2022MNRAS.516..942C", "2022MNRAS.516.3924V", "2022MNRAS.517.1557R", "2022arXiv220309051H", "2022arXiv220311818H", "2023A&A...670A.100L", "2023A&A...671A.102E", "2023A&A...676A..26G", "2023MNRAS.518.6293G", "2023MNRAS.519.4651H", "2023arXiv230310077C", "2024A&A...684A..23T", "2024A&A...686A..98R", "2024AJ....168...12S", "2024MNRAS.527.2624P", "2024MNRAS.527.5792L", "2024MNRAS.527.6506B", "2024MNRAS.528.2326S", "2024MNRAS.531.3551L", "2024MNRAS.tmp.1483J", "2024MNRAS.tmp.1484K", "2024arXiv240108750G", "2024arXiv240414103F", "2024arXiv240510908N", "2024arXiv240605179Z", "2024arXiv240611987P", "2024arXiv240614613N" ]
[ "astronomy" ]
37
[ "galaxies: elliptical and lenticular", "cD", "galaxies: formation", "galaxies: fundamental parameters", "galaxies: spiral", "galaxies: statistics", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stw1495", "10.48550/arXiv.1607.01096" ]
1607
1607.01096_arXiv.txt
At the fundamental level galaxies are multi-component systems \citep[see for example][]{Buta2010}, consisting of at least a spheroid and/or disc. This is most obvious in the S\'ersic index -- colour plane where the single component Sd and elliptical galaxies occupy distinct peaks with composite galaxies (S0abc) scattered between and around these peaks \citep[see e.g.,][]{Driver2006,Cameron2009,Kelvin2012,Lange2015}. These components have very different characteristics with spheroids typically having a featureless appearance and being pressure supported. Discs on the other hand have features such as spiral arms and are rotationally supported. Furthermore bulges are made up of redder stars with moderate to high metallicities and a high $\alpha$-element abundance, while discs are made of younger, bluer stars with lower metallicities and typically are dust and gas rich. Spheroids are older, showing little to no star formation and are typically dust and gas depleted \citep[see for example the review by][]{RobertsHaynes1994}. The simplest explanation for these stark differences is that spheroids and discs form via two distinct mechanisms over two distinct eras \citep{Cook2009,Driver2013}, i.e. a dynamically \mbox{``hot mode''} (spheroid formation) and ``cold mode'' (disc formation) evolution.\\ Traditionally the relative prominence of a bulge component is taken into account when classifying galaxies onto the Hubble sequence (see \citealt{Hubble1926}, and later revisions by e.g., \citealt{vdBergh1976,Kormendy2012}), however studying global properties of galaxies by Hubble type could be misleading. For example, numerous evolution mechanisms have been proposed to explain the morphological diversity seen at $z=0$, such as a (initial) major dissipative event, gas accretion, adiabatic contraction, major and minor mergers and secular processes \citep[see e.g.,][]{Hopkins2010,Trujillo2011,L'Huillier2012,Cheung2013,Sachdeva2015}. Each of these processes potentially acts to modify the prominence of the bulge, disc or other components. This indicates that galaxy components likely follow distinct formation pathways and structure effectively encodes the formation history. Therefore, to study galaxy evolution \bd{} decomposition is critical. While the number of studies of large samples which employ \bd{} decomposition to explore the nature of galaxies and their components is growing, the analysis is challenging \citep[see e.g.,][]{Allen2006,Gadotti2009,Simard2011,Bruce2012,Bruce2014,Lang2014,Tasca2014,Meert2015,Salo2015}. This is because multi-component fitting is notoriously difficult, especially when trying to automate it for large samples. Nevertheless a number of publicly available codes have now been created to allow \bd{} decomposition, such as \textsc{gim2d} \citep{Simard1998}, \textsc{budda} \citep{deSouza2004}, \textsc{galfit3} \citep{Peng2010} and \textsc{imfit} \citep{Erwin2015}. Each code has advantages and disadvantages \citep[see][for example for further discussion]{Erwin2015}, here we elect to use \galfit 3 because of its ability to manage nearby objects, its computational reliability, and its speed. Many studies that fit 2-component S\'ersic light profiles restrict the S\'ersic index to n=1 for the disc and in some cases n=4 for the bulge \citep[e.g.~][]{Simard2011,Bruce2012,Lackner2012,Meert2015}. This reduces the number of free parameters and ensures the fitting process is more robust but it restricts the possible interpretations of the fitting outcomes, e.g. classical and pseudo bulges can not be differentiated this way. A number of studies now show that the S\'ersic index of discs and spheroids (be they pure or component) vary smoothly with mass and luminosity or due to dust or galaxy type \citep[see e.g.][]{Graham2003,Gadotti2009,Kelvin2012,Graham2013,Pastrav2013a,Pastrav2013b}. Hence studies where the S\'ersic index of the bulge or disc components are fixed may be overly restrictive. Furthermore, to correctly trace a galaxy's formation history a full decomposition of all of its components would be ideal \citep[e.g., the Spitzer Survey of Stellar Structure in Galaxies, S$^4$G,][]{Salo2015}. However, this is only viable for very nearby galaxies where all the components can be clearly resolved and hence for relatively small samples (S$^4$G is the largest study to date extending to 2352 galaxies for which a number have been fit with more than 2 components). To compare to galaxies at different epochs going beyond a simple bulge and disc decomposition is difficult \citep{Gadotti2008}. There are two reasons, however, why two components might be sufficient, (i) the majority of stellar mass resides in the bulge and disc components for most galaxies, and (ii) some components may simply represent minor perturbations to the underlying disc (e.g. bars, pseudo-bulges). Such perturbations should arguably be considered secondary rather than primary evolutionary markers. Here we adopt the stance that bulge and disc components arise from two primary formation pathways (i.e. hot and cold mode evolution, respectively), and that additional components form in secondary formation pathways (i.e., tidal interactions, disc instabilities and perturbations). The likely primary pathways are: monolithic collapse followed by major mergers, which can produce elliptical galaxies by destroying and rearranging any structure previously present in a galaxy, resulting in a smooth light profile \citep{Toomre1977}; and minor mergers and continued gas inflow, which can form or re-grow a disc around a pre-existing spheroid, resulting in a galaxy with two distinct components \citep[see e.g.,][]{Steinmetz2002,Kannappan2009,Wei2010}. A key question worth asking is whether two generic components (spheroids and discs) really can explain the diversity seen, i.e., how many fundamental building blocks and structures are required to adequately reproduce the observed galaxy population? As most of the stellar mass is contained within the bulge and disc how important are tertiary features like bars? Furthermore how many different physical origins do the various spheroids and discs have? Are elliptical galaxies simply naked bulges and are bulges related to high-redshift compact galaxies \citep[e.g.][]{Graham2015,Berg2014}? Are the discs of early-types, late-types and irregulars indistinguishable? \\ We believe that the stellar mass -- half-light size (hereafter \msr{}) relation is a key scaling relation allowing us to address these questions for the following reasons: \begin{itemize} \item The size of a galaxy is related to its specific angular momentum making the mass and size of a galaxy fundamental observables of conserved quantities \citep[e.g.,][]{Romanowsky2012}. \item The simple assumption that angular momentum is conserved during the initial collapse of the dark matter halo links the angular momentum and mass of a galaxy with its dark matter halo \citep{Fall1980,Dalcanton1997,Mo1998}. \item Hydrodynamical simulations now produce galaxies with realistic sizes and direct comparisons (at different epochs) are possible to study formation and evolution histories of galaxies \citep[see for example the Evolution and Assembly of Galaxies and their Environments simulation suite, \eagle,][]{Schaye2015,Crain2015}. \item We can empirically measure and trace the masses and sizes of galaxies and their components over a range of redshifts and in different environments (e.g., with HST as well as high-redshift ground-based surveys and soon with Euclid and WFIRST). \end{itemize} The \msr{} relation therefore represents the next critical diagnostic for galaxy evolution studies beyond simple mass functions \citep[see e.g.,][]{Bouwens2004,Wel2014,Holwerda2015,Shibuya2015}, enabling us to trace angular momentum build-up and the emergence of the component nature of galaxies while connecting observations to simulations.\\ Recent studies comparing the \msr{} relation of low and high-redshift are already yielding interesting results. For example, at high-redshift galaxies might look disc-like or elliptical\slash spheroidal but their physical properties are unlike any discs or ellipticals in the local Universe \citep[see e.g.~][]{Bruce2012,Buitrago2013,Mortlock2013}. Galaxies at high redshifts are typically more irregular with thick slab-like disc structures and clumpy star-forming regions \citep{Wisnioski2012}. In addition, they can be very compact but massive. In some cases, at redshift $\sim$2 they are a factor of up to 6 times smaller in size than galaxies of the same mass today \citep{Daddi2005,Trujillo2007,Buitrago2008,vanDokkum2008,vanDokkum2010,Weinzirl2011}.\\ In this paper we aim to provide a reliable low redshift benchmark of the \msr{} relation for bulges, discs and spheroids. The \bd decomposition sample is derived from a set of galaxies for which detailed morphological information is available (see \citealt{Moffett2015}). Section \ref{sec:data} describes the data and sample selection, Sections \ref{sec:2comp} and \ref{sec:compmass} describe the set up of our \bd decomposition catalogue and component mass estimates. In Section \ref{sec:msr} we present the \msr{} relations for bulges, spheroidal and disc galaxies and discuss the association of components with their possible parent populations. We then compare our distributions to the \eagle\ simulation in Section \ref{sec:sims} followed by a comparison of our low redshift \msr{} relation with recent high redshift data from Ultra Deep Survey (UDS) region within the Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey \citep[CANDELS,][]{Grogin2011,Koekemoer2011} in Section \ref{sec:highz}. Finally in Section \ref{sec:SnC} we present our summary and conclusions.\\ Throughout this paper we use data derived from the Galaxy And Mass Assembly (GAMA) survey \citep{Driver2011,Driver2016,Liske2014} with stellar masses derived from \cite{Taylor2011}, sizes derived from S\'ersic profile fitting using \textsc{sigma} \citep{Kelvin2012}, and for a cosmology given by: $\Lambda{}$ Cold Dark Matter universe with $\Omega_{\mathrm{m}} = 0.3,\; \Omega_{\Lambda}=0.7,\; H_0 = 70 \mathrm{km s}^{-1}\;\mathrm{Mpc}^{-1}$.
\label{sec:SnC} We have presented our bulge-disc decomposition catalogue for 7506 galaxies from the GAMA survey in the redshift range of $0.002<z<0.06$ (Sec.~\ref{sec:2comp}). To overcome the limitations of the LM minimisation algorithm used in \galfit{}, which can get trapped in local minima (especially for 2-component fits), we repeatedly fit our galaxy sample with varying starting points to map out the parameter space. For the single component galaxies we use a set of 33 combinations of the starting parameters, and for the 2-component fits we use a set of 88 combinations. We implement a screening process to prune bad fits and determine the final fitting values and errors from the median of the acceptable fits. We use the 16$^{th}$ and 84$^{th}$ percentile of the remaining output parameter distribution to determine the error on the median model combined with a 10\% error floor. Through this strategy we reduce our catastrophic failure rate from $\sim 20\%$ to $\sim 5\%$.\\ We then presented the \msr{} relations of our sample by Hubble type and component with the component masses based on an estimation from the bulge and disc colours. Next we explored the association of the bulge and disc components with either the Sd-Irr or elliptical \msr{} relation. We find that S(B)ab-S(B)cd galaxies likely consist of a disc plus a pseudo-bulge. Considering that a pseudo-bulge is a perturbation of the disc we decide that our late-type 2-component systems are best represented by a single component S\'ersic fit. Thus we associate elliptical, early-type bulges and LBS for the spheroid \msr{} relation and Sd-Irr, single component fit S(B)ab-S(B)cd galaxies and early-type discs for the final disc \msr{} relation, which we provide as a definitive low redshift benchmark:\\ \\ $R_e=5.141 \mathrm{\left(\frac{\mathcal{M}_*} {10^{10}\mathcal{M}_{\odot}}\right)}^{0.274}$ for discs and,\\ \\ $R_e=2.063 \mathrm{\left(\frac{\mathcal{M}_*} {10^{10}\mathcal{M}_{\odot}}\right)}^{0.263}$ for spheroids.\\ However, we caution the reader that the spheroid relation is heavily dominated by low mass galaxies. If a comparison to high mass spheroids is needed then the high mass elliptical \msr{} relation (see Table \ref{table:compfits}) should be used in lieu of a curved spheroid relation.\\ Next we used our local disc and spheroid \msr{} distributions to compare to data from the \eagle\ simulation. We find a qualitatively good agreement between the sizes of the \eagle\ data and our \msr{} relations. This is not surprising as the sizes in \eagle\ were calibrated using the \citep{Shen2003} \msr{} distribution. Hence the variance is of more interest as this has not been explicitly matched between the observed and simulated data. Comparing the scatter of the observed and simulated data we find that in almost all cases the simulated data has a smaller scatter which is unexpected, considering that the dark matter spin distribution is known to be fairly broad and we would expect the sizes and angular momentum distributions to have similarly broad distributions. Since we are comparing half-light sizes to half-mass sizes and components versus active and passive galaxies, we caution the reader to not over-interpret this comparisons. Instead we would like to highlight the potential of using the mass-size plane to compare observational and simulated data. Finally, we compare our local \msr{} relations to high redshift data from the CANDELS-UDS field. We concentrate on available data in a redshift range of $1<z\leq 1.5$ with available single and 2-component fits \citep{Mortlock2013,MB2015}. We generally find that low-mass high redshift galaxies agree better with the local \msr{} distributions than high-mass high redshift galaxies. Furthermore, high redshift systems, with the exception of disc components, more closely follow that of our local spheroid relation. The high redshift discs on the other hand follow the local disc \msr{} relation, albeit offset to slightly smaller size. We interpret this as evidence for spheroid formation at high redshift and propose that further disc formation and/or growth does not occur until later times.
16
7
1607.01096
We perform automated bulge + disc decomposition on a sample of ∼7500 galaxies from the Galaxy And Mass Assembly (GAMA) survey in the redshift range of 0.002 &lt; z &lt; 0.06 using Structural Investigation of Galaxies via Model Analysis, a wrapper around GALFIT3. To achieve robust profile measurements, we use a novel approach of repeatedly fitting the galaxies, varying the input parameters to sample a large fraction of the input parameter space. Using this method, we reduce the catastrophic failure rate significantly and verify the confidence in the fit independently of χ<SUP>2</SUP>. Additionally, using the median of the final fitting values and the 16th and 84th percentile produces more realistic error estimates than those provided by GALFIT, which are known to be underestimated. We use the results of our decompositions to analyse the stellar mass - half-light radius relations of bulges, discs and spheroids. We further investigate the association of components with a parent disc or elliptical relation to provide definite z = 0 disc and spheroid M_star - R_e relations. We conclude by comparing our local disc and spheroid M_star - R_e to simulated data from EAGLE and high-redshift data from Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey-Ultra Deep Survey. We show the potential of using the M_star - R_e relation to study galaxy evolution in both cases but caution that for a fair comparison, all data sets need to be processed and analysed in the same manner.
false
[ "Deep Extragalactic", "Survey-Ultra Deep Survey", "Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey-Ultra Deep Survey", "simulated data", "disc decomposition", "elliptical relation", "discs", "galaxy evolution", "∼7500 galaxies", "Mass Assembly", "spheroids", "lt", "GALFIT3", "Model Analysis", "χ", "Structural Investigation", "the input parameter space", "Cosmic Assembly Near-infrared", "z", "EAGLE" ]
11.941066
6.22697
189
12436535
[ "Breed, M.", "Venter, C.", "Harding, A. K." ]
2016arXiv160706480B
[ "Very-high energy emission from pulsars" ]
2
[ "-", "-", "-" ]
[ "2016heas.confE..40V", "2017ASSL..446...61D" ]
[ "astronomy" ]
6
[ "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.48550/arXiv.1607.06480" ]
1607
1607.06480_arXiv.txt
Since the launch in June 2008 of {\it Fermi} LAT \cite{Atwood2009}, a high-energy (HE) satellite measuring $\gamma$-rays in the range 20 MeV$-$300 GeV, two pulsar catalogues (1PC, \cite{Abdo2010}; 2PC, \cite{Abdo2013}) discussing the light curve and spectral properties of 117 pulsars have been released. The vast majority of the {\it Fermi}-detected pulsars display exponentially cutoff spectra with cutoffs around a few GeV. These spectra are believed to be due to curvature radiation (CR), which is assumed to be the dominating emission process in the GeV band (see Section~\ref{subsection:radiationmechanisms}). \subsection{Standard pulsar emission models}\label{subsection:emissionmodels} There exist several physical radiation models that can be used to study HE emission from pulsars. These include the polar cap (PC; \cite{Daugherty1982}), slot gap (SG; \cite{Arons1983}), outer gap (OG; \cite{Cheng1986b}), and the pair-starved polar cap (PSPC; \cite{Harding2005}) models, which can be distinguished from each other based on the different assumptions of the geometry and location of the `gap regions'. The `gap region' is where particle acceleration takes place due to an unscreened, rotation-induced $E$-field parallel to the local $B$-field, as well as subsequent emission by these particles. In PC models emission from HE particles is assumed to originate close to the neutron star (NS) surface. These particles are accelerated by large $E$-fields near the magnetic poles (known as the magnetic PCs) only up to a few stellar radii. In SG models, the radiation comes from narrow gaps close to the last open field lines (the field lines that are tangent to the light cylinder where the corotation speed equals the speed of light $c$), with the gaps extending from the NS surface up to high altitudes. In the OG model, the gap region extends from the null-charge surface, where the Goldreich-Julian charge density is zero \cite{Goldreich1969} up to high altitudes, also close to the last open field lines. The PSPC model involves a gap region that extends from the NS surface to the light cylinder over the full open volume \cite{Harding2005}, since the potential is unscreened in this case, so that there are not enough pairs to fully screen the $E$-field. \subsection{Radiation and pair creation processes}\label{subsection:radiationmechanisms} To explain HE emission in the standard models, one has to take detailed particle transport and radiation mechanisms into account. These mechanisms include CR, synchrotron radiation (SR), and inverse Compton scattering (ICS). CR occurs whenever charged particles are constrained to move along curved paths, e.g., along curved $B$-field lines (e.g., \cite{Harding1981}), therefore involving a change in their longitudinal kinetic energy. When the emitted CR photon energy and the local $B$-field are high enough, magnetic pair production may occur (where an HE photon converts into an electron-positron pair, $e^\pm$), leading to a cascade of $e^\pm$ pairs which may screen the parallel $E$-field outside the gaps. The pair cascade is characterized by the so-called multiplicity, i.e., the number of pairs spawned by a single primary. The pairs may radiate SR if they have velocity components perpendicular to the local $B$-field so that this process involves a change in the particles' transverse kinetic energy. Also, ICS occurs due to the relativistic particles which upscatter soft photons (e.g., originating at a heated PC), which results in the ``boosting'' of the photon energies up to very high energies. ICS photons may also be converted into $e^\pm$ pairs. Two-photon pair creation ($\gamma\gamma$-absorption) may also occur, in particular in OG models. Cyclotron emission combined with subsequent ICS has also been considered by \cite{Lyutikov2013} to explain the broadband spectrum of the Crab pulsar. \subsection{Historic perspective of VHE spectral modelling} \begin{figure}[t]\centering \includegraphics[width=22.3pc,height=13pc]{figure1.eps}\hspace{0.5cm} \caption{\label{Romani96} Early prediction of the phase-averaged spectrum for the Vela pulsar. The solid lines represent the spectral components for CR, SR, and the thermal surface flux (kT). The dashed curve represents the TeV pulsed spectral component associated with the ICS of SR of the primary $e^\pm$ (SSC). Adapted from \cite{Romani1996}.} \end{figure} Early modelling, assuming the standard OG model, predicted spectral components in the VHE regime when estimating the ICS of primary electrons on SR or soft photons. This resulted in a natural bump around a few TeV (involving $\sim 10$ TeV particles) in the extreme Klein-Nishina limit as seen in figure~\ref{Romani96} and \ref{Hirotani01}. However, these components may not survive up to the light cylinder and beyond \cite{Cheng1986b,Romani1996,Hirotani2001a}, since $\gamma\gamma$ pair creation leads to absorption of the TeV $\gamma$-ray flux. \begin{figure}[t]\centering \includegraphics[width=20pc,height=15pc]{figure2.eps}\hspace{0.5cm} \caption{\label{Hirotani01} Expected TeV spectra for three bright pulsars including the Crab (solid lines), PSR B0656$+$14 (dashed lines), and PSR B1509$-$58 (dotted lines). The thick and thin curves represent inclination angles (between the spin and magnetic axis) of $\alpha=30^\circ$ and $45^\circ$, respectively \cite{Hirotani2001a}.} \end{figure} Other studies assumed CR to be the dominant radiation mechanism producing $\gamma$-ray emission when performing spectral modelling and found spectral cutoffs of up to 50 GeV. For example, \cite{Bulik2000} modelled the cutoffs of millisecond pulsars (MSPs). These are pulsars which possess relatively low $B$-fields and short periods. Their model assumed a static dipole $B$-field and a PC geometry, and predicted CR from the primary electrons that are released from the PC. Their predicted CR spectral component cut off at $\sim100$~GeV. The CR photons may undergo pair production in the intense low-altitude $B$-fields, and the newly formed electron-positron secondaries will emit SR in the optical and X-ray band (see \cite{Harding2002}). Therefore, they concluded that the HE CR from MSPs occurred in an energy band that was above the detection range of satellite detectors like \textit{Energetic Gamma-Ray Experiment Telescope (EGRET)} and below that of ground-based Cherenkov detectors such as the \textit{High Energy Stereoscopic System (H.E.S.S. I)}. Later studies investigated the X-ray and $\gamma$-ray spectrum of rotation-powered MSPs using a PSPC model \cite{Harding2005}, and found CR cutoffs of $\sim10-50$~GeV (see \cite{Frackowiak2005,Venter2005}). Optical to $\gamma$-ray spectra were also modelled by \cite{Harding2008} assuming an SG accelerator and a retarded vacuum dipole (RVD) $B$-field, for the Crab pulsar. They found spectral cutoffs of up to a few GeV. Another study modelled the phase-resolved spectra of the Crab pulsar using the OG and SG models, and found HE cutoffs of up to $\sim25$~GeV \cite{Hirotani2008}. Cutoffs around $\sim10$~GeV were found for the OG model using the RVD $B$-field \cite{Tang2008}.
The abovementioned detections of the VHE pulsed emission from pulsars and the explanation thereof implies that this emission may yield strong constraints on $\gamma$-ray radiation mechanisms, the location of acceleration regions, and the $B$-field structure. There is thus an urgent need for refinements and extensions of standard pulsar emission models and radiation mechanisms, including more realistic $B$-fields. Some examples are discussed above, but there are many more. SG model refinements include photon-photon pair production attenuation within the model and also more realistic $B$-fields such as dissipative magnetospheric solutions. One could also model the emission pulse profiles as a function of energy. Ground-based Cherenkov telescopes are now searching for more examples of VHE pulsars. New pulsar models will assist us in predicting the level of VHE emission expected from them, which would be very important for the upcoming \textit{CTA}. The low threshold energy of \textit{CTA} will provide an overlap with the {\it Fermi} energy range and will help to discriminate between CR and a potentially new spectral component. \ack This work is based on research supported by the National Research Foundation (NRF) of South Africa (Grant Numbers 90822, 93278, and 99072). Any opinions, findings, and conclusions or recommendations expressed are that of the authors, and the NRF accepts no liability whatsoever in this regard. A.K.H.\ acknowledges the support from the NASA Astrophysics Theory Program.
16
7
1607.06480
The vast majority of pulsars detected by the Fermi Large Area Telescope (LAT) display exponentially cutoff spectra with cutoffs falling in a narrow band around a few GeV. Early spectral modelling predicted spectral cutoffs at energies of up to 100 GeV, assuming curvature radiation. It was therefore not expected that pulsars would be visible in the very-high energy (VHE) regime (&gt;100 GeV). The VERITAS announcement of the detection of pulsed emission from the Crab pulsar at energies up to 400 GeV (and now up to 1.5 TeV as detected by MAGIC) therefore raised important questions about our understanding of the electrodynamics and local environment of pulsars. H.E.S.S. has now detected pulsed emission from the Vela pulsar down to tens of GeV, making this the second pulsar detected by a ground-based Cherenkov telescope. Deep upper limits have also been obtained by VERITAS and MAGIC for the Geminga pulsar. We will review the latest developments in VHE pulsar science, including an overview of the latest observations, refinements, and extensions to radiation models and magnetic field structures, and the implementation of new radiation mechanisms. This will assist us in understanding the VHE emission detected from the Crab pulsar, and predicting the level of VHE emission expected from other pulsars, which is very important for the upcoming CTA.
false
[ "new radiation mechanisms", "radiation models", "curvature radiation", "magnetic field structures", "spectral cutoffs", "other pulsars", "pulsars", "VHE pulsar science", "VHE pulsar", "VHE emission", "cutoffs", "pulsed emission", "VHE", "energies", "local environment", "important questions", "CTA", "the Crab pulsar", "gt;100", "MAGIC" ]
5.534463
4.28026
69
12407656
[ "Kim, Woong-Tae", "Moon, Sanghyuk" ]
2016ApJ...829...45K
[ "Equilibrium Sequences and Gravitational Instability of Rotating Isothermal Rings" ]
3
[ "Center for the Exploration of the Origin of the universe (CEOU), Department of Physics &amp; Astronomy, Seoul National University, Seoul 151-742, Korea ; Center for Theoretical Physics (CTP), Seoul National University, Seoul 151-742, Korea", "Center for the Exploration of the Origin of the universe (CEOU), Department of Physics &amp; Astronomy, Seoul National University, Seoul 151-742, Korea" ]
[ "2018MNRAS.475...63H", "2019ApJS..241...24M", "2021MNRAS.505.4310C" ]
[ "astronomy" ]
3
[ "galaxies: ISM", "galaxies: kinematics and dynamics", "galaxies: nuclei", "galaxies: structure", "instabilities", "stars: formation", "Astrophysics - Astrophysics of Galaxies" ]
[ "1953mtp..book.....M", "1964ApJ...140.1056O", "1964ApJ...140.1067O", "1965MNRAS.130...97G", "1967ApJ...147..334C", "1968ApJ...151.1075O", "1971ApJ...167..425B", "1981PThPh..65.1870E", "1985A&A...150..327C", "1986ApJS...61..479H", "1986ApJS...61..609B", "1986ApJS...62..461H", "1987PThPh..77..635N", "1988ApJ...324...60B", "1988MNRAS.231...97G", "1990MNRAS.245..614L", "1990Natur.345..679S", "1991A&A...244..257G", "1992ApJ...388..392I", "1994ApJ...425L..73E", "1994PASJ...46..243M", "1994tisp.book.....G", "1995ApJ...454..623K", "1996AJ....111.1861B", "1996ARA&A..34..155B", "1996FCPh...17...95B", "1997ApJS..108..471A", "1997NewA....2....1C", "1999ApJ...527...86C", "1999acfp.book.....L", "2000AJ....120.1289B", "2000AN....321..363C", "2001ApJ...559...70K", "2001MNRAS.323..663R", "2002AJ....123.1411B", "2002ApJ...570..132K", "2005ApJ...632..217S", "2006A&A...448..489K", "2006ApJ...646..213K", "2006ApJ...649..181S", "2006MNRAS.371.1087A", "2008AJ....135..479B", "2008ApJS..174..337M", "2008gady.book.....B", "2009ApJ...698..715S", "2010A&A...518L..59S", "2010MNRAS.402.2462C", "2010PASA...27...56R", "2011A&A...529A..45V", "2011Ap&SS.334....1H", "2011ApJ...736..129H", "2011ApJ...739..104M", "2011MNRAS.412.2396F", "2012A&A...543A..61B", "2012ApJ...747...60K", "2012ApJ...751..124K", "2012ApJ...758...14K", "2012ApJ...761..131K", "2013A&A...551A..81V", "2013ApJ...769..100S", "2014Ap&SS.353..191H", "2014ApJ...789...68K", "2014ApJ...792...47S", "2015ApJ...806...39O", "2015ApJ...806..150L", "2015ApJ...809...33K", "2015MNRAS.450...53H", "2016MNRAS.455...51H" ]
[ "10.3847/0004-637X/829/1/45", "10.48550/arXiv.1607.03570" ]
1607
1607.03570_arXiv.txt
\label{s:intro} Nuclear rings in barred-spiral galaxies often exhibit strong activities of star formation (e.g, \citealt{but96,ken97,kna06,maz08,san10,maz11,hsi11,van11, oni15}). They are mostly circular, with ellipticity of $e\sim0-0.4$. They are thought to form due to nonlinear interactions of gas with an underlying non-axisymmetric stellar bar potential (e.g., \citealt{com85,but86,shl90,kna95,com01,com10}). Recent hydrodynamic simulations show that the inflowing gas driven inward by the bar torque tends to gather at the location of centrifugal barrier, well inside the inner Lindblad resonance, where the centrifugal force on the gas balances the external gravity \citep{kim12b,kim12c,ks12,li15}. This predicts that nuclear rings are smaller in size in galaxies with stronger bars and/or lower pattern speeds, overall consistent with observational results of \citet{com10}. One of important issues regarding nuclear rings is what determines the star formation rate (SFR) in them. Observations indicate that the ring SFRs vary widely in the range of $\sim0.1$--$10\Mrate$ from galaxy to galaxy, with a smaller value corresponding to a more strongly barred galaxy \citep{maz08,com10}, although the total gas mass in each ring is almost constant at $\sim(1$--$6) \times 10^8\Msun$ (e.g., \citealt{but00,ben02,she05,sch06}). By analyzing photometric H$\alpha$ data of 22 nuclear rings, \citet{maz08} found that about a half of their sample possesses an azimuthal age gradient of young star clusters in such a way that older ones are located systematically farther away from the contact points between a ring and dust lanes, while other rings do not show a noticeable age gradient (see also, e.g., \citealt{bok08,ryd10,bra12}). To explain the spatial distributions of ages of young star clusters, \citet{bok08} proposed two scenarios of star formation: the ``popcorn'' model in which star formation takes place in dense clumps randomly distributed along a nuclear ring, and the ``pearls on a string'' model where star formation occurs preferentially near the contact points. Since star clusters age as they orbit about the galaxy center, the pearls-on-a-string model naturally explains the presence of an azimuthal age gradient, while clusters with different ages are well mixed in the popcorn model (see also, e.g., \citealt{ryd01,ryd10,all06,san10,van13}). The most important factor that determines the dominating type of star formation appears to be the mass inflow rate $\dot{M}$ to the ring along the dust lanes \citep{seo13,seo14}. When $\dot{M}$ is less than a critical value $\dot{M}_{c}$, all inflowing gas can be consumed at the contact points, and star formation occurs in the pearls-on-a-string fashion. When $\dot{M}> \dot{M}_{c}$, on the other hand, the inflowing gas overflows the contact points and is transferred into other parts of the ring, resulting in popcorn-style star formation when it becomes gravitationally unstable. \citet{seo13} found numerically $\dot{M}_c\sim 1\Mrate$ for typical nuclear rings, although it depends rather sensitively on the gas sound speed as well as the ring size. The above consideration implicitly assumes that nuclear rings undergoing star formation in the pearls-on-a-string manner are gravitationally stable, while those with popcorn-type star formation are globally unstable. However, this has yet to be tested theoretically. Although several authors studied gravitational instability of ring-like systems (e.g., \citealt{goo88,elm94,chr97,had11}), it is difficult to apply their results directly to nuclear rings because of the approximations made in these studies. For example, \citet{goo88} analyzed a linear stability of shearing accretion rings (or tori) to gravitational perturbations, but their models were limited to incompressible gas without any motion along the vertical direction parallel to the rotation axis (see also \citealt{luy90,and97}). For magnetized compressible rings, \citet{elm94} showed that a ring with density larger than $0.6\kappa^2/G$ is gravitationally unstable, with $\kappa$ and $G$ referring to the epicycle frequency and the gravitational constant, respectively. However, this result was based on the local approximation that treated the ring as a thin uniform cylinder without considering its internal structure. On the other hand, \citet{had11} and \citet{had14} analyzed stability of polytropic rings with index $n=1.5$ by solving the linearized equations as an initial value rather than eigenvalue problem, and found several unstable modes with the azimuthal mode number $m\leq 4$. \citet{chr97} instead ran two-dimensional nonlinear simulations of galaxy rings using the equations integrated along the vertical direction. Using an adiabatic equation of state, they found that massive slender rings are highly unstable to gravitating modes with $m$ as large as 18. However, these linear or nonlinear initial-value approaches did not search all unstable modes systematically as functions of $m$ and rotation frequency $\Omega_0$. Is a ring with given physical quantities (such as mass, size, sound speed, rotation speed) gravitationally stable or not? What is the most dominant mode if it is unstable? How fast does it grow? To address these questions, we in this paper perform a linear stability analysis of nuclear rings, assuming that they are slender and isothermal. We will find full dispersion relations of gravitationally unstable modes as well as the critical angular frequencies for stability. We will then apply the results to observed nuclear rings to check the presence or absence of an azimuthal age gradient of young star clusters is really consistent with stability properties of the rings. We will also run three-dimensional numerical simulations and compare the results with those of our linear stability analysis. Stability analysis of any system requires to set up its initial equilibrium a priori. Due to their complicated geometry, finding equilibrium configurations of isothermal rings is a non-trivial task. In a pioneering work, \citet{ost64b} treated the effects of rotation and the curvature as perturbing forces to otherwise non-rotating infinite cylinders, and obtained approximate expressions for density distributions of polytropic or isothermal rings in axisymmetric equilibrium. To determine the equilibrium structure of a slowly-rotating, spheroid-like body, \citet{ost68} developed a self-consistent field (SCF) method that solves the Poisson equation as well as the equation for steady equilibrium, alternatively and iteratively. \citet{eri81} used a similar iteration method to find a ring-like equilibrium sequence of incompressible bodies as a function of $\Omega_0$. \citet{hac86a,hac86b} extended the original SCF method of \citet{ost68} to make it work even for rapidly-rotating, ring-like polytropes in two or three dimensions. In this paper, we shall modify the SCF technique of \citet{hac86a} to find equilibrium sequences of rigidly-rotating isothermal bodies. This will allow us to explore the effects of compressibility on the internal structures of rings. The remainder of this paper is organized as follows. In Section \ref{s:eql}, we describe our SCF method used to construct isothermal bodies in steady equilibrium. In Section \ref{s:seq}, we present the equilibrium sequences of rigidly-rotating isothermal objects, together with test results for incompressible bodies and Bonner-Ebert spheres. We will also show that the density profiles of slender rings can well be approximated by those of infinite isothermal cylinders. In Section \ref{s:GI}, we perform a linear stability analysis of slender isothermal rings to obtain the dispersion relations as well as the critical angular frequencies, and present the results of numerical simulations. In Section \ref{s:sum}, we summarize and conclude this work with applications to observed nuclear rings.
\label{s:sum} \subsection{Summary} Nuclear rings at centers of barred galaxies exhibit strong star formation activities. They are thought to undergo gravitational instability when sufficiently massive. To study their equilibrium properties and stability to gravitational perturbations, we approximate nuclear rings as isothermal objects. We modify the SCF method of \citet{hac86a} to make it suitable for an isothermal equation of state, and construct equilibrium sequences of rigidly-rotating, self-gravitating, isothermal bodies. A steady equilibrium is uniquely specified by two dimensionless parameters: $\alpha$ and $\hatrB$ (see Eqs.~[\ref{e:alp}] and [\ref{e:rB}]). The former is the measure of the thermal energy relative to gravitational potential energy of an equilibrium body, while the latter corresponds to the ellipticity for spheroid-like configurations or the thickness for ring-like configurations. We take a convention that $\hatrB$ is positive (or negative) for spheroid-like (or ring-like) objects. To test our SCF method, we first apply it to the case of rotating incompressible bodies, and confirm that our method is able to reproduce the Maclaurin spheroid sequence when $0.158 \leq \hatrB \leq1$. With improved resolution, our method gives more accurate results than those obtained by \citet{eri81} and \citet{hac86a} for the concave hamburger sequence with $0 \lesssim \hatrB < 0.158$ . Our method also successfully reproduces isothermal Bonnor-Ebert spheres, with larger $\alpha$ corresponding to a higher degree of central density concentration. We then use our SCF method to obtain the density distributions of rotating isothermal equilibria on the meridional plane, as illustrated in Figure \ref{f:isocontour}. We calculate the dependence on $\hatrB$ of various dimensionless quantities such as the rotational angular frequency $\hatOmegas$, the total mass $\hatM$, the mean density $\hatavgrho$, the total kinetic energy $\hatT$, and the gravitational potential energy $\hatW$. These values are tabulated in Table \ref{t:iso} and given graphically in Figure \ref{f:isoRB}. We find that an equilibrium density profile is more centrally concentrated for smaller $\alpha$. Unlike the incompressible bodies, not all values of $\hatrB$ result in an isothermal equilibrium configuration. Spheroid-like equilibria exist only for $\hatrBmax \leq \hatrB \leq 1$, while ring-like (or hamburger-like) configurations are possible only for $-1< \hatrB<\hatrBmin$: otherwise, the centrifugal potential is too large to form gravitationally bound objects. The critical $\hatrB$ values are found to be $\hatrBmax=0.27$, 0.51, and 0.59, and $\hatrBmin=0.13$, $-0.14$, and $-0.40$ for $\alpha=1$, $0.1$, and $0.01$, respectively. In general, $\hatOmegas$ is a decreasing function of $|\hatrB|$. This is naturally expected for spheroid-like configurations since faster rotation leads to a more flattened equilibrium. As $\hatrB$ approaches $-1$, on the other hand, ring-like configurations becomes less massive and thus requires weaker centrifugal force to balance self-gravity. Due to stronger central concentration, $\hatOmegas$, $\hatM$, $\hatavgrho$, $\hatT$, and $|\hatW|$ all become smaller as $\alpha$ decreases. For $\alpha< 0.1$, $\hatM$ and $\hatW$ are insensitive to $\hatrB\gtrsim0.6$ for spheroid-like equilibria since the density in the outer parts becomes vanishingly small. For a given value of the normalized angular momentum $j$, the normalized angular frequency $\omega_s$ becomes smaller with decreasing $\alpha$, although the energy ratio $T/|W|$ is insensitive to $\alpha$. \begin{deluxetable*}{rcccccccc}[!t] \tablecaption{Properties of Observed Nuclear Rings\label{t:obsring}} \tablewidth{0pt} % \tablehead{ & \colhead{$R_1$} & & \colhead{$v_{\rm rot}$} & \colhead{$M_g$} & & & & \\ \colhead{Galaxy} & \colhead{(kpc)} & \colhead{$e$} & \colhead{(km s$^{-1}$)} & \colhead{$(10^7\Msun)$} & \colhead{Age Grad.} & \colhead{$\alpha$} & \colhead{$\hatOmega$} & \colhead{Ref.} \\ \colhead{(1)} & \colhead{(2)} & \colhead{(3)} & \colhead{(4)} & \colhead{(5)} & \colhead{(6)} & \colhead{(7)} & \colhead{(8)} & \colhead{(9)} } \startdata NGC 473 & 1.69 & 0.06 & 125 & 40 & Yes & 1.69E$-2$ & 1.71 & \\ NGC 613 & 0.40 & 0.26 & 115 & 40 & ? & 3.93E$-3$ & 0.76 & \\ NGC 1097 & 0.97 & 0.32 & 220 & 140 & No & 2.70E$-3$ & 1.20 & (a),(b) \\ NGC 1300 & 0.40 & 0.15 & 155 & 40 & No & 3.98E$-3$ & 1.03 & \\ NGC 1343 & 1.97 & 0.30 & 80 & 40 & Yes & 1.92E$-2$ & 1.17 & \\ NGC 1530 & 1.20 & 0.80 & 180 & 40 & Yes & 9.30E$-3$ & 1.82 & \\ NGC 2903 & 0.16 & 0.32 & 60 & 35 & ? & 1.78E$-3$ & 0.27 & (c) \\ NGC 3351 & 0.15 & 0.11 & 120 & 31 & ? & 1.89E$-3$ & 0.55 & (c) \\ NGC 4303 & 0.35 & 0.11 & 90 & 42 & No & 3.32E$-3$ & 0.54 & \\ NGC 4314 & 0.56 & 0.31 & 160 & 21 & Yes & 1.04E$-2$ & 1.71 & (d),(e) \\ NGC 4321 & 0.87 & 0.32 & 170 & 51 & Yes & 6.64E$-3$ & 1.45 & \\ NGC 5248 & 0.65 & 0.20 & 150 & 42 & Yes & 6.13E$-3$ & 1.23 & \\ NGC 5728 & 1.10 & 0.23 & 180 & 40 & Yes & 1.09E$-2$ & 1.97 & \\ NGC 5905 & 0.39 & 0.14 & 150 & 40 & ? & 3.88E$-3$ & 0.98 & \\ NGC 5953 & 1.00 & 0.43 & 150 & 40 & ? & 9.50E$-3$ & 1.54 & \\ NGC 6951 & 0.56 & 0.17 & 160 & 40 & Yes & 5.56E$-3$ & 1.25 & \\ NGC 7552 & 0.34 & 0.15 & 150 & 40 & No & 3.38E$-3$ & 0.92 & (f) \\ NGC 7716 & 1.20 & 0.04 & 150 & 40 & No & 1.20E$-2$ & 1.72 & \\ NGC IC14 & 0.68 & 0.36 & 204 & 40 & Yes & 6.57E$-3$ & 1.74 & \enddata \tablecomments{ Columns (2) and (3) give the semi-major axis and ellipticity of nuclear rings adopted from \citet{com10}. Column (4) is the rotational velocity adopted from \citet{maz08} or from the references given in Column (9). Column (5) is the total gas mass in the ring from \citet{she05} or from references in Column (9); we take $M_g = 4\times 10^8\Msun$ if no information is available. Column (6) cites the age distribution: ``Yes'' and ``No'' for the presence and absence of an azimuthal age gradient, respectively, and ``?'' for uncertain cases, adopted from \citet{all06}, \citet{maz08}, \citet{san10}, and \citet{bra12}. Columns (7) and (8) give $\alpha$ and $\hatOmegas$ calculated by Equations \eqref{e:obs_alp} and \eqref{e:obs_Omg}. Column (9) is the references for $v_{\rm rot}$ or $M$: (a) \citet{oni15}; (b) \citet{hsi11}; (c) \citet{maz11}; (d) \citet{ gar91}; (e) \citet{ben96}; (f) \citet{bra12}.} \end{deluxetable*} The density distribution of finite slender rings obtained by our SCF method for $\hatrB \lesssim -0.6$ is found to be well approximated by Equation \eqref{e:rhosr}, which is also the solution for static, isothermal cylinders of infinite extent. This indicates that the rotation as well as geometrical curvature effect are insignificant in determining an equilibrium for rings with the major axis $R_0$ much longer than the minor axis $\eta_0$. The equilibrium angular frequency for isothermal slender rings with $\alpha\gtrsim 0.1$ is well described by Equation \eqref{e:Ostinc} applicable to truncated incompressible rings \citep{ost64b}. To explore gravitational instability of nuclear rings, we calculate the growth rates of nonaxisymmetric modes with azimuthal mode number $m$ by assuming that the rings are slender with $\eta_0/R_0=0.1$, and that perturbations are independent of the polar angle $\lambda$ in the meridional plane. In the absence of rotation, the resulting dispersion relations are the same as those of axisymmetric modes for an infinite isothermal cylinder studied by \citet{nag87} if $m/R_0$ is taken equal to the wavenumber in the direction along the cylinder (see Fig.~\ref{f:rdisp_all}). Only large-scale modes can be gravitationally unstable, and the unstable range of $m$ as well as the maximum growth rate increase with decreasing $\alpha$. Rotation tends to stabilize the gravitational instability, reducing both the growth rates and the unstable ranges of $m$. The instability is completely suppressed when $\hatOmega$ exceeds the critical value that is relatively constant at $\sim 0.7$ for $\alpha\gtrsim0.01$ and increases rapidly with decreasing $\alpha$ (see Fig.~\ref{f:crit_cal}). The simple estimates of the critical angular frequencies from the Toomre condition as well as the local dispersion relation are smaller than the results of our full stability analysis for $\alpha \lesssim 5\times 10^{-3}$ due to underestimation of self-gravity at the ring centers. Shear turns out to be unimportant for the gravitational instability of rings as long as they are slender. \subsection{Discussion}\label{s:dis} \citet{maz08} analyzed photometric data of a sample of nuclear rings to estimate the ages of H$\alpha$-emitting star clusters and found that about a half of their sample contains an age gradient of star clusters along the azimuthal direction. Since nuclear rings with age gradient are thought to be gravitationally stable and form stars preferentially at the contact points, it is interesting to apply the results of our linear stability analysis to the observed rings to tell whether they are really stable. Table \ref{t:obsring} lists the properties of 19 observed nuclear rings in galaxies with a noticeable bar, compiled from the literature where the information on the presence/absence of an age gradient is available.\footnote{Most of the galaxies listed in Table \ref{t:obsring} except for NGC 1097, NGC 2903, NGC 3351, and NGC 7752 are adopted from \citet{maz08}: NGC 1097 is from \citet{san10}, NGC 2903 and NGC 3351 from \citet{maz11}, NGC 4321 from \citet{all06}, and NGC 7752 from \citet{bra12}.} Column (1) lists each galaxy name. Columns (2) and (3) give the semi-major axis and ellipticity of nuclear rings adopted from \citet{com10}. Column (4) lists the rotational velocity $v_{\rm rot}$ adopted from \citet{maz08} or from the references given in Column (9). Column (5) lists the total gas mass $M_g$ in the ring from \citet{she05} or the references in Column (9) only for the galaxies with available data; we otherwise take $M_g=4\times10^8\Msun$ as a reference value. Column (6) indicates the presence or absence of an azimuthal age gradient of star clusters adopted from \citet{maz08}, \citet{all06}, \citet{san10}, and \citet{bra12}: a question mark is used when it is difficult to characterize the age distribution. Columns (7) and (8) give $\alpha$ and $\hatOmega$ calculated by \be\label{e:obs_alp} \alpha = 0.01 \left(\frac{c_s}{10\kms}\right)^2 \left(\frac{M_g}{4\times10^8\Msun}\right)^{-1} \left(\frac{R_0}{1\kpc}\right), \ee and \be\label{e:obs_Omg} \hatOmega = 2.1 \left(\frac{v_{\rm rot}}{200\kms}\right) \left(\frac{M_g}{4\times10^8\Msun}\right)^{-0.5} \left(\frac{R_0}{1\kpc}\right)^{0.5}, \ee after taking $R_0=R_1(1-e^2)^{1/4}$ corresponding to the geometric means of the major and minor axes of eccentric rings, $\cs=10\kms$, and $\eta_0=R_0/10$. We replace $\rho_c$ with $\avgrho$ since the ring central density is difficult to constrain observationally. In Figure \ref{f:crit_obs}, we plot $\hatOmega$ against $\alpha$ for the rings listed in Table \ref{t:obsring} using various symbols, with numbers indicating galaxy names. Overall, rings with larger $\alpha$ tend to have larger $\hatOmega$. Blue circles represent rings with an azimuthal age gradient, while red diamonds are for those with no age gradient. Rings for which the age distribution cannot be judged are marked by star symbols. It is apparent that all rings with an azimuthal age gradient are located at the stable regime, while all rings with no age gradient, except NGC 7716, correspond to unstable configurations. These results are consistent with two modes of star formation proposed by \citet{bok08}, such that rings sufficiently massive or rotating sufficiently slowly form stars in the popcorn style caused by gravitational instability, and thus do not show an apparent age gradient. On the other hand, star formation in stable rings may occur preferentially at the contact points to exhibit an azimuthal age gradient of star clusters like pearls on a string. The ring models we have considered so far ignored the effects of magnetic fields that are pervasive in galaxies (e.g., \citealt{bec96,fle11}). In spiral galaxies, the presence of toroidal magnetic fields is known to play a destabilizing role in forming giant clouds inside spiral arms where tension forces from bent field lines resist the stabilizing effect of the Coriolis force (e.g., \citealt{bal88,kim01,kim02,kim06}). In addition, magnetic fields are likely to reduce the degree of central density concentration by exerting pressure forces. It will be interesting to study how magnetic fields embedded in nuclear rings change the critical angular frequencies for gravitational instability compared to those of unmagnetized rings. \begin{figure} \includegraphics[angle=0, width=8.5cm]{fig16.pdf} \caption{Distributions of $\alpha$ and $\hatOmega$ of the observed nuclear rings listed in Table \ref{t:obsring}. Blue circles and red diamonds represent rings with and without an azimuthal age gradient, respectively, while rings with uncertain age distributions are indicated by star symbols. \label{f:crit_obs}} \end{figure}
16
7
1607.03570
Nuclear rings at the centers of barred galaxies exhibit strong star formation activities. They are thought to undergo gravitational instability when they are sufficiently massive. We approximate them as rigidly rotating isothermal objects and investigate their gravitational instability. Using a self-consistent field method, we first construct their equilibrium sequences specified by two parameters: α corresponding to the thermal energy relative to gravitational potential energy, and {\widehat{R}}<SUB>{{B</SUB>}} measuring the ellipticity or ring thickness. Unlike in the incompressible case, not all values of {\widehat{R}}<SUB>{{B</SUB>}} yield an isothermal equilibrium, and the range of {\widehat{R}}<SUB>{{B</SUB>}} for such equilibria shrinks with decreasing α. The density distributions in the meridional plane are steeper for smaller α, and well approximated by those of infinite cylinders for slender rings. We also calculate the dispersion relations of non-axisymmetric modes in rigidly rotating slender rings with angular frequency Ω<SUB>0</SUB> and central density {ρ }<SUB>c</SUB>. Rings with smaller α are found more unstable with a larger unstable range of the azimuthal mode number. The instability is completely suppressed by rotation when Ω<SUB>0</SUB> exceeds the critical value. The critical angular frequency is found to be almost constant at ∼ 0.7{(G{ρ }<SUB>c</SUB>)}<SUP>1/2</SUP> for α ≳ 0.01 and increases rapidly for smaller α. We apply our results to a sample of observed star-forming rings and confirm that rings without a noticeable azimuthal age gradient of young star clusters are indeed gravitationally unstable.
false
[ "slender rings", "ring thickness", "smaller α", "Rings", "rings", "Nuclear rings", "young star clusters", "strong star formation activities", "gravitational potential energy", "α", "such equilibria", "gravitational instability", "a larger unstable range", "non-axisymmetric modes", "observed star-forming rings", "∼ 0.7{(G{ρ", "the azimuthal mode number", "SUB", "isothermal objects", "a noticeable azimuthal age gradient" ]
10.325019
7.020678
133
12464733
[ "Salzano, Vincenzo", "Mota, David F.", "Dabrowski, Mariusz P.", "Capozziello, Salvatore" ]
2016JCAP...10..033S
[ "No need for dark matter in galaxy clusters within Galileon theory" ]
15
[ "Institute of Physics, University of Szczecin, Wielkopolska 15, 70-451 Szczecin, Poland", "Institute of Theoretical Astrophysics, University of Oslo, 0315 Oslo, Norway", "Institute of Physics, University of Szczecin, Wielkopolska 15, 70-451 Szczecin, Poland; National Centre for Nuclear Research, Andrzeja Sołtana 7, 05-400 Otwock, Poland ; Copernicus Center for Interdisciplinary Studies, Sławkowska 17, 31-016 Kraków, Poland;", "Dipartimento di Fisica ``E. Pancini'', Università degli Studi di Napoli ``Federico II'', and INFN—Sezione di Napoli, Complesso Universitario di Monte S. Angelo, Via Cinthia, Edificio N, 80126 Napoli, Italy ; Gran Sasso Science Institute (INFN), Viale F. Crispi, 7, I-67100, L'Aquila, Italy;" ]
[ "2018ApJ...856..177C", "2018EPJC...78..447C", "2018IJMPD..2730003M", "2018PDU....19..137H", "2019arXiv190708010V", "2020EPJC...80...24G", "2020PhRvD.102j4063G", "2021EPJP..136..235G", "2021MNRAS.506..595P", "2021arXiv210512582S", "2022EPJC...82..652B", "2022MNRAS.511.1878L", "2022PhRvD.105l4024M", "2022arXiv220406048K", "2023PhRvD.107d4008B" ]
[ "astronomy" ]
3
[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1972PhLB...39..393V", "1974IJTP...10..363H", "1978A&A....70..677C", "1996ApJ...462..563N", "2005MNRAS.358..601C", "2006ApJ...638L..51M", "2007PhRvD..75f3501M", "2007PhRvD..75f3508B", "2008GReGr..40..357C", "2009PhRvD..79f4036N", "2009PhRvD..79h4003D", "2009arXiv0912.0914L", "2010PhLB..694..198L", "2010PhRvD..82l4006G", "2010arXiv1001.0061R", "2011PhR...509..167C", "2011PhRvD..84f4039D", "2011arXiv1110.3193L", "2012ApJS..199...25P", "2012PhR...513....1C", "2012PhRvL.109j1802L", "2013CQGra..30u4006D", "2013LRR....16....6A", "2013PhR...530...87W", "2014ApJ...794..136D", "2014IJMPD..2343010G", "2014JCAP...08..059B", "2014LRR....17....4W", "2014PhRvD..89f4046Z", "2014arXiv1401.0046M", "2015ApJ...806....4M", "2015CQGra..32x3001B", "2015JCAP...02..018G", "2015JHEP...12..055A", "2015PhR...568....1J", "2015PhRvD..91l4066K", "2015PhRvD..92l4045S", "2015PhRvL.114u1101G", "2015PhRvL.115o9901A", "2015PhRvL.115t1101S", "2015Sci...349..849H", "2016A&A...594A..13P", "2016A&A...594A..14P", "2016ARNPS..66...95J", "2016JCAP...07..019S", "2016PDU....12...56B", "2016PhRvL.116f1102A", "2016PhRvL.116x1103A" ]
[ "10.1088/1475-7516/2016/10/033", "10.48550/arXiv.1607.02606" ]
1607
1607.02606_arXiv.txt
\label{sec:Introduction} Despite the astonishing successes which have been accomplished very recently \citep{GW1,GW2}, following the high-level-accuracy cosmological observations collected so far \citep{PlanckCosmo,PlanckMod}, and meanwhile waiting for newest and greatly improved projects to be launched and fully operative (SKA\footnote{\url{https://www.skatelescope.org/}} and \textit{Euclid}\footnote{\url{http://sci.esa.int/euclid/}} \citep{Laureijs09,Laureijs11,Refregier10,Amendola13} among others), the full validity of general relativity (GR) over all the possible astrophysical and cosmological ranges is still under debate (see \citep{BeyondLCDM} and reference therein). This is mainly due to the enigmatic nature of two of its main pillars which are nowadays believed to rule the dynamics of the entire universe and of everything is inside it: Dark Matter (DM) and Dark Energy (DE). It is indisputable that no candidate for DM has yet been detected; some break in the Standard Model of Physics might be possible, as stated in \citep{LHC1,LHC2,LHC3}, but the results have not been statistically confirmed yet, and they might be only loosely related to the DM problem. Also, we are basically unaware about the nature of DE \citep{DE1,DE2} and willing for more decisive statistical evidences from observations \citep{DE3}. What could be good or bad news is that we have plenty of alternatives to GR, at least for what concerns the interpretation and the solution of DE-related problems \citep{BeyondLCDM,Clifton12,clifton}. And all can be bounded, and many times disproved, by requiring GR as a limit at such scales where we know it works well as, for example, Solar System scales \citep{Will14}. The family-theories that generally pass this test are based on some sort of \textit{screening mechanism}: the new introduced degrees of freedom (whatever they are, in the form of a new scalar field, or as geometrical effects) are generally suppressed at small scales, so that any deviation from GR cannot be detected \citep{Hamilton15,shaw,gano}, but stay active at larger cosmological scales, where they can mimic DE, with no ad-hoc assumptions and in a theoretical apparatus possibly much more general than GR. Even in this case, anyway, we have a generous selection of mechanisms of various origin which might produce such screenings (see \citep{Joyce15,bour,Berti15} for a general review). Here, we will focus on a well-defined scenario: a Galileon-type alternative model, which spontaneously breaks the underlying Vainshtein screening mechanism. Galileons are a particular class of scalar field invariant under galilean shift symmetry and which, despite of leading to higher-derivative field theories, still have second order equations of motions \citep{Nicolis09,Deffayet09,Deffayet11,Deffayet13}. Generally, they are able to pass Solar System scale test by means of a the so-called Vainshtein mechanism \citep{Vainshtein72} due to the particular way the kinetic contributions of the Galileon fields are defined, with first or second order derivatives becoming important at a certain scale. Such a theory has been widely studied; actually, we will focus on a particular version of it, recently proposed in \citep{KoyamaSakstein2015}. This model is very interesting because it naturally contains a breaking of the Vainshtein mechanism at some (to be defined) scale, which basically implies that there should be some breaking of GR at some point, and this should be detectable \cite{mauro}. One possible way to test such a hypothesis has been proposed in \citep{SaksteinPRL15,SaksteinPRD15} in a stellar physics context. A generalized version of it has been discussed in \citep{Sakstein2016} and tested with clusters of galaxies, using both gravitational lensing and X-ray observations. Here, we will try to face a very different approach: such a model has been considered in all the above cases only as a general alternative to GR, in particular, as a good candidate for DE, but no suggestion about a possible role played in substituting DM has been given. In general, most of these alternative theories are suggested in order to explain the dynamics of the universe on cosmological scales without requiring the presence of an exotic fluid like DE. In the specific case of Galileon models, they have been studied in details as cosmological background and, as such, influencing the formation of gravitational structures. But no direct connection with a possible role in mimicking DM has been defined. Actually, even theoretically, such possibility was in some way prohibited by the same screening mechanism. But with the new model developed in \citep{KoyamaSakstein2015,Sakstein2016} we have a natural breaking of such mechanism so that one could naturally ask: \textit{what if the mechanism were broken at some astrophysical scale(s) and, thus, the Galileon model might play the role of DM in some way?} In order to give an answer to this question, we will approach the problem in the following way. First, we have to identify the kind of observations which might result to be helpful for our purposes. We found it in the convergence map that can be derived from the analysis of strong and weak lensing events from clusters of galaxies. Actually, clusters of galaxies are the best gravitational objects in order to combine both phenomena: they constitute massive and large enough foreground lenses able to produce detectable distortions of background sources both near the center (where the depth of the gravitational potential creates strong lensing events) and in the outer regions (where the cluster mass produces the statistical distortions of background galaxies known as weak lensing). Finally, the convergence map will result to be, basically, the two-dimensional projected mass distribution of the clusters that can be reconstructed by taking into account all these lensing events. Clusters of galaxies are also characterised by another dynamical property: the large amount of intra-cluster gas (the largest matter component, if we do not consider dark matter) is heated up by the deep gravitational potential and emits in the X-ray band. X-ray observations of hot gas in clusters of galaxies provide us with another tool to have information about the internal mass distribution of a cluster. As it is well known, there might be internal astrophysical phenomena which can locally perturb the gas, or even larger-scale events like merging from smaller sub-structures, which can heat it up in addition to the pure gravitational attraction, leading to not properly correct mass estimations. Mass reconstruction from lensing, on the other hand, is much more accurate because it is insensitive to local dynamics, and can reproduce the true mass distribution due to pure gravitational attraction quite well. Actually, we will not be interested directly on these aspects in this work, but more on the fact that, through X-ray observations, the gas density distribution of the cluster can be measured with very high confidence. Such gas density will be then used as input to calculate the convergence map from lensing in the context of the Galileon model we have discussed above. The important point to be stressed is that the gas will be the only contribution to the cluster mass we will consider; no additional DM will be assumed here. The paper is organized as follows: in section~(\ref{sec:Model}) we describe the theoretical apparatus at the base of our analysis; in section~(\ref{sec:Data}) we describe what kind of data we have considered and how we implement them; in section~(\ref{sec:Results}) we discuss the results obtained and their implications; finally, in section~(\ref{sec:Conclusions}), we summarize all our analysis.
\label{sec:Conclusions} In this work we have tried to answer the question: is it possible to explain clusters of galaxies lensing observations without resorting to DM? This is clearly impossible in the classical context of GR where, actually, DM is vital and needed in order to explain observations. But it might be possible in some alternative models of gravity. Most of those have been introduced to explain large scale effects of DE; then, they need some screening mechanism to turn off the new forces/degrees of freedom at such scales where GR should be recovered. The same screening mechanisms are the main obstacle to elect many alternative models as DM sources too, because the suppression of the new degrees of freedom generally happens at astrophysical scales. A safety alternative come from models like that described in \citep{KoyamaSakstein2015,Sakstein2016}, a modified Galileon-type model where the screening mechanism can be broken; in such a case, distinctive signatures of the new theory should be observable. And, we add, they might play the role of DM. We have used the above model and applied it to a sample of $18$ clusters of galaxies observed by the spatial survey \textit{CLASH}. We have joined two different observational probes: mass reconstruction from lensing events (both strong and weak lensing); and X-ray observations of the hot intra-cluster gas. We have used the latter to derive the gas mass profile of such clusters; and we have used this information trying to reproduce the former observations. The key requirement in our analysis is that the mass of the clusters is made only by gas, no DM is considered at all. Anyway, one should not forget that baryonic matter in clusters is also made of galaxies. Unfortunately, we have no data available for the sample we considered in this sense. Including galaxies, of course, might affect our results: in particular, galaxies could play a larger role in inner regions, with a consequent ``re-normalization'' of the mass profile, and possible influence on the theoretical parameters. Results show that when the screening mechanism is broken, then the Galileon model can be used to match DE at large cosmological scales, and DM at smaller ones. In particular, it results that the Galileon model is even more statistically favourable than the GR to match lensing observations. Far from us to state that this ends the DM question, or that such model is ``the'' model which can win GR. But we think this analysis is useful to state that much more attention should be paid to GR alternative and competing models \cite{rev}, because they can be as much successful as GR. And it could help theoreticians to understand what is the right path to follow in order to build/discover the true underlying gravity theory behind our universe.
16
7
1607.02606
Modified gravity theories with a screening mechanism have acquired much interest recently in the quest for a viable alternative to General Relativity on cosmological scales, given their intrinsic property of being able to pass Solar System scale tests and, at the same time, to possibly drive universe acceleration on much larger scales. Here, we explore the possibility that the same screening mechanism, or its breaking at a certain astrophysical scale, might be responsible of those gravitational effects which, in the context of general relativity, are generally attributed to Dark Matter. We consider a recently proposed extension of covariant Galileon models in the so-called ``beyond Horndeski'' scenario, where a breaking of the Vainshtein mechanism is possible and, thus, some peculiar observational signatures should be detectable and make it distinguishable from general relativity. We apply this model to a sample of clusters of galaxies observed under the CLASH survey, using both new data from gravitational lensing events and archival data from X-ray intra-cluster hot gas observations. In particular, we use the latter to model the gas density, and then use it as the only ingredient in the matter clusters' budget to calculate the expected lensing convergence map. Results show that, in the context of this extended Galileon, the assumption of having only gas and no Dark Matter at all in the clusters is able to match observations. We also obtain narrow and very interesting bounds on the parameters which characterize this model. In particular, we find that, at least for one of them, the general relativity limit is excluded at 2σ confidence level, thus making this model clearly statistically different and competitive with respect to general relativity.
false
[ "general relativity", "Solar System scale tests", "cosmological scales", "X-ray intra-cluster hot gas observations", "Dark Matter", "gravitational lensing events", "observations", "covariant Galileon models", "clusters", "Solar System", "archival data", "universe acceleration", "the general relativity limit", "a certain astrophysical scale", "General Relativity", "much larger scales", "the expected lensing convergence map", "no Dark Matter", "much interest", "respect" ]
10.338744
1.359662
-1
2073358
[ "Ruppin, F.", "Adam, R.", "Comis, B.", "Ade, P.", "André, P.", "Arnaud, M.", "Beelen, A.", "Benoît, A.", "Bideaud, A.", "Billot, N.", "Bourrion, O.", "Calvo, M.", "Catalano, A.", "Coiffard, G.", "D'Addabbo, A.", "De Petris, M.", "Désert, F. -X.", "Doyle, S.", "Goupy, J.", "Kramer, C.", "Leclercq, S.", "Macías-Pérez, J. F.", "Mauskopf, P.", "Mayet, F.", "Monfardini, A.", "Pajot, F.", "Pascale, E.", "Perotto, L.", "Pisano, G.", "Pointecouteau, E.", "Ponthieu, N.", "Pratt, G. W.", "Revéret, V.", "Ritacco, A.", "Rodriguez, L.", "Romero, C.", "Schuster, K.", "Sievers, A.", "Triqueneaux, S.", "Tucker, C.", "Zylka, R." ]
2017A&A...597A.110R
[ "Non-parametric deprojection of NIKA SZ observations: Pressure distribution in the Planck-discovered cluster PSZ1 G045.85+57.71" ]
36
[ "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France; Laboratoire Lagrange, Université Côte d'Azur, Observatoire de la Côte d'Azur, CNRS, Bd de l'Observatoire, CS 34229, 06304, Nice Cedex 4, France", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Astronomy Instrumentation Group, University of Cardiff, UK", "Laboratoire AIM, CEA/IRFU, CNRS/INSU, Université Paris Diderot, CEA-Saclay, 91191, Gif-Sur-Yvette, France", "Laboratoire AIM, CEA/IRFU, CNRS/INSU, Université Paris Diderot, CEA-Saclay, 91191, Gif-Sur-Yvette, France", "Institut d'Astrophysique Spatiale (IAS), CNRS and Université Paris Sud, Orsay, France", "Institut Néel, CNRS and Université Grenoble Alpes, France", "Institut Néel, CNRS and Université Grenoble Alpes, France", "Institut de RadioAstronomie Millimétrique (IRAM), Granada, Spain", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Institut Néel, CNRS and Université Grenoble Alpes, France", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Institut de RadioAstronomie Millimétrique (IRAM), Grenoble, France", "Institut Néel, CNRS and Université Grenoble Alpes, France; Dipartimento di Fisica, Sapienza Università di Roma, Piazzale Aldo Moro 5, 00185, Roma, Italy", "Dipartimento di Fisica, Sapienza Università di Roma, Piazzale Aldo Moro 5, 00185, Roma, Italy", "Institut de Planétologie et d'Astrophysique de Grenoble (IPAG), CNRS and Université Grenoble Alpes, France", "Astronomy Instrumentation Group, University of Cardiff, UK", "Institut Néel, CNRS and Université Grenoble Alpes, France", "Institut de RadioAstronomie Millimétrique (IRAM), Granada, Spain", "Institut de RadioAstronomie Millimétrique (IRAM), Grenoble, France", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Astronomy Instrumentation Group, University of Cardiff, UK; School of Earth and Space Exploration and Department of Physics, Arizona State University, Tempe, AZ, 85287, USA", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Institut Néel, CNRS and Université Grenoble Alpes, France", "Institut d'Astrophysique Spatiale (IAS), CNRS and Université Paris Sud, Orsay, France", "Astronomy Instrumentation Group, University of Cardiff, UK", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Astronomy Instrumentation Group, University of Cardiff, UK", "Université de Toulouse, UPS-OMP, Institut de Recherche en Astrophysique et Planétologie (IRAP), 31028, Toulouse, France; CNRS, IRAP, 9 avenue Colonel Roche, BP 44346, 31028, Toulouse Cedex 4, France", "Institut de Planétologie et d'Astrophysique de Grenoble (IPAG), CNRS and Université Grenoble Alpes, France", "Laboratoire AIM, CEA/IRFU, CNRS/INSU, Université Paris Diderot, CEA-Saclay, 91191, Gif-Sur-Yvette, France", "Laboratoire AIM, CEA/IRFU, CNRS/INSU, Université Paris Diderot, CEA-Saclay, 91191, Gif-Sur-Yvette, France", "Laboratoire de Physique Subatomique et de Cosmologie, Université Grenoble Alpes, CNRS/IN2P3, 53 avenue des Martyrs, Grenoble, France", "Laboratoire AIM, CEA/IRFU, CNRS/INSU, Université Paris Diderot, CEA-Saclay, 91191, Gif-Sur-Yvette, France", "Institut de RadioAstronomie Millimétrique (IRAM), Grenoble, France", "Institut de RadioAstronomie Millimétrique (IRAM), Grenoble, France", "Institut de RadioAstronomie Millimétrique (IRAM), Granada, Spain", "Institut Néel, CNRS and Université Grenoble Alpes, France", "Astronomy Instrumentation Group, University of Cardiff, UK", "Institut de RadioAstronomie Millimétrique (IRAM), Grenoble, France" ]
[ "2016sf2a.conf..439D", "2017A&A...606A..64A", "2017AN....338..305P", "2017arXiv170506064P", "2017arXiv170901255M", "2017ehep.confE..42M", "2018A&A...612A..39R", "2018A&A...614A.118A", "2018A&A...615A..18R", "2018A&A...615A.112R", "2018MNRAS.481..749S", "2018arXiv180810817P", "2019A&A...621A..41G", "2019A&A...631A..21R", "2019A&A...632A..22C", "2019BAAS...51c.124M", "2019MNRAS.486.2116P", "2019MNRAS.487.4037D", "2019SSRv..215...17M", "2019SSRv..215...25P", "2020A&A...634A.113A", "2020A&A...637A..71P", "2020A&A...642A..60C", "2020A&A...644A..93K", "2020ApJ...893...74R", "2020EPJWC.22800009E", "2020EPJWC.22800016M", "2020EPJWC.22800017M", "2020EPJWC.22800020P", "2021ApJ...918...43R", "2022A&A...661A..65F", "2023A&A...671A..28M", "2023A&A...678A.197M", "2023OJAp....6E...9K", "2024A&A...684A..18A", "2024arXiv240300909D" ]
[ "astronomy" ]
39
[ "galaxies: clusters: intracluster medium", "instrumentation: high angular resolution", "cosmology: observations", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1051/0004-6361/201629405", "10.48550/arXiv.1607.07679" ]
1607
1607.07679_arXiv.txt
\label{sec:Introduction} Galaxy clusters are the ultimate manifestation of the hierarchical structure formation process in the standard cosmological model, and as such, they are sensitive to both the matter content and expansion history of the Universe in which they form. Clusters are thus potentially powerful tools to infer cosmological parameters. In particular, counting clusters as a function of their mass and redshift \citep[\emph{e.g.,}][]{seh11,Planck_SZ_cosmo2015,deh16} brings constraints on the cosmological parameters that are complementary to those derived with other probes such as type Ia supernovae \citep[\emph{e.g.,}][]{rie07}, the CMB temperature and polarization angular power spectra \citep[\emph{e.g.,}][]{Planck_param}, or baryonic acoustic oscillations \citep[\emph{e.g.,}][]{and14}.\\ About 85\% of the total mass in galaxy clusters is from dark matter. The principal baryonic component is found in the hot, ionized, X-ray emitting gas of the intracluster medium (ICM), containing about 12\% of the total mass. The remaining baryonic mass is found in the stellar population. Cluster masses can be inferred from several independent observables. The velocity dispersion of the galaxies \citep[\emph{e.g.,}][]{biv06,sif16}, various X-ray properties such as temperature or luminosity \citep[\emph{e.g.,}][]{SVM_prof,pra09}, or the lensing distortions of background galaxies \citep[\emph{e.g.,}][]{app14,ume14,hoe15} can be related to the underlying total mass. Another observational probe of interest is the thermal Sunyaev-Zel'dovich effect \citep[tSZ;][]{sun72}, which is due to the inverse Compton scattering of cosmic microwave background (CMB) photons with high-energy electrons of the ICM. As this effect is directly proportional to the thermal energy contained in the ICM, it is expected to provide a low scatter mass proxy for galaxy clusters \citep[\emph{e.g.,}][]{das04,nag07}. Furthermore, as the tSZ effect is a CMB spectral distortion, it does not suffer from cosmological dimming. This observable is therefore a powerful probe to estimate both galaxy cluster total mass and baryonic content distribution up to high redshift.\\ The \planck\ satellite, the South Pole Telescope (SPT), and the Atacama Cosmology Telescope (ACT) surveys have used tSZ observations to discover and characterize large galaxy cluster samples \citep[\emph{e.g.,}][]{Planck_cata2,SPTcluster,ACT_cluster}. In addition, individual observations of known clusters have been obtained with a number of instruments, such as APEX-SZ, CARMA, SZA, BOLOCAM, and AMIs \citep[\emph{e.g.,}][]{Apexsz,Carmasz,SZA,Bolocamsz,AMI_followup}. However, their relatively low angular resolution ($> 1$ arcmin) restricts the tSZ characterization of the ICM to low redshift \citep{pla10, NonparamPressure,bon12,Planck_pressure_prof,SayersPointSource}, as a combination with higher resolution X--ray observations is needed to map clusters at both large and small scales \citep[\emph{e.g.,}][]{Planck_pressure_prof,eck13}.\\ In combination with local data, high angular resolution tSZ observations at intermediate to high-redshift ($z > 0.5$) have a number of different applications. They can be used to study the evolution of structural properties such as cluster pressure profiles and their scatter. Furthermore, they provide new insights and constraints on scaling properties such as the relation between the integrated Compton parameter and the cluster total mass and its scatter. High angular resolution tSZ observations can also be used to characterize the two-dimensional (2D) pressure distribution within the ICM. This information is essential for understanding cluster formation physics and performing precise cosmological analysis with the cluster population.\\ Cluster growth and evolution is characterized by complex astrophysical phenomena, including deviation from equilibrium and generation of turbulence due to merging events and feedback from active galactic nuclei. While stochastic, the frequency of these events evolves with time and increases at high redshift. They are the prime cause of scatter and deviations from self-similarity in the scaling relations that are used to link observables to mass in cosmological analyses (\emph{e.g.,} \citealt{mergers_bias,MUSIC_scaling}). Of particular importance is the clarification of the physical origin of this normalization and scatter in the scaling relations, rendering the use of galaxy clusters for cosmological application more robust.\\ The New IRAM KIDs Array \citep[NIKA;][]{NIKA_cam1,NIKA_elec2,cal13} was a dual-band continuum camera operated at the Institut de Radio Astronomie Millimetrique (IRAM) 30 m telescope between 2010 and 2015. It was one of the very few tSZ instruments with sub-arcminute resolution. Other examples include the Goddard-IRAM Superconducting 2-Millimeter Observer \citep[GISMO;][]{sta08} and the Multiplexed SQUID TES array at Ninety Gigahertz \citep[MUSTANG;][]{kor11}. NIKA was the only dual-band sub-arcminute instrument \citep{NIKA_calib} that observed the tSZ effect simultaneously at 150 and 260~GHz with an angular resolution of 18.2 and 12.0~arcsec, respectively. Furthering the characterization of galaxy cluster pressure profiles that has been initiated by arcminute resolution instruments at low redshift, NIKA has now mapped the pressure distribution in a number of galaxy clusters at intermediate and high redshift \citep[see][]{RXJ1347NIKA,CLJ1227NIKA,MACSJ1424NIKA,MACSJ0717NIKA}.\\ In this paper we detail the NIKA observations of the \planck-discovered cluster \psz\ at $z=0.61$. A key result is the first non-parametric measurement with high statistical precision of the pressure profile of a distant cluster at an angular resolution $\sim 20$ arcsec, extending to much higher redshift pioneering non-parametric pressure profile measurements at low resolution \citep{NonparamPressure}. \citealt{NonparamPressure} have applied the deprojection method presented in \citealt{NonparamNord} to the APEX-SZ data \citep{NonparamApex} of the nearby cluster Abell 2204 ($z=0.15$). They have shown that a non-parametric modeling of the gas pressure profile can be obtained. Previous works have shown that deprojection methods can be used to probe the ICM of clusters from simulations \citep{NonparamSimu1,NonparamSimu2,NonparamSimu3}.\\ The work presented in this paper is a pilot study for the forthcoming SZ observations (see \citealt{ProceedingMoriond}) with NIKA2 (see \citealt{ProceedingSPIE}). The combination with \planck\ data allows the determination of the non-parametric pressure profile out to scales of $\gtrsim 3$ Mpc, substantially improving the constraints on the spherically integrated Compton parameter. Using the deprojected gas density profile from \xmm, we reconstruct the thermodynamic properties of the ICM without making use of X-ray spectroscopic information. This result illustrates the excellent synergy between tSZ and X-ray observations of similar angular resolution, and serves as a pilot study for combining tSZ data to measure the gas pressure with short X-ray observations to measure the gas density.\\ This paper is organized as follows. The NIKA observations of \psz\ at the IRAM 30-meter telescope and the raw data processing are explained in Sect. 2. Ancillary data, previous tSZ observations, point source contamination, and XMM-Newton data reduction, are described in Sect. 3. The modelization of the ICM and the method to estimate the cluster total mass are presented in Sect. 4. We also discuss the characterization of the cluster ellipticity and its impact on the mass estimation. In Sect. 5 a non-parametric multiprobe analysis is performed to extract the radial pressure profile and obtain the ICM thermodynamic properties. The conclusions and NIKA2 perspectives are discussed in Sect. 6. Throughout this study we assume a flat $\Lambda$CDM cosmology following the latest \planck\ results \citep{Planck_param}: H$_0 = 67.8$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm m} = 0.308$, and $\Omega_\Lambda = 0.692$. Within this framework, at the cluster redshift, one arcsec corresponds to 6.93 kpc. \begin{figure*}[h] \centering \includegraphics[height=6.8cm]{PSZ1G045_map_2mm_for_paper.pdf} \includegraphics[height=6.8cm]{PSZ1G045_map_1mm_for_paper.pdf} \caption{{\footnotesize NIKA tSZ surface brightness maps at 150 GHz (left) and 260 GHz (right). The significance of the measured signal is given by the black contours starting at $3\sigma$ with $1\sigma$ spacing. The maps are smoothed with an additional 10~arcsec Gaussian filter for display purposes and the NIKA beam FWHMs are represented as white disks in the bottom left-hand corner of the maps. The white crosses indicate the X-ray center. Note that we use the original maps (without additional smoothing) in the following analysis.}} \label{fig:brightness_map} \end{figure*}
\label{sec:conclusions} The \planck\ tSZ-discovered cluster \psz\ has been observed simultaneously at 150 and 260~GHz by the NIKA camera. A 4.35 hour observation allowed a detailed mapping at 18.2 arcsec angular resolution of the tSZ signal at 150~GHz. The cluster was also observed in the X-ray band by the \xmm\ satellite. We performed the first non-parametric pressure profile deprojection from resolved tSZ observations of a \planck-discovered cluster at an intermediate redshift ($z=0.61$). The MCMC procedure, which was developed to deproject the cluster pressure profile, uses the NIKA tSZ surface brightness map and the \planck\ Compton parameter map jointly to constrain the cluster pressure distribution from its core up to $5R_{500}$. The resulting pressure profile does not deviate significantly from the standard gNFW model. The combination of both NIKA and \planck\ data brings strong constraints on the pressure profile slope at each scale, and allows a significant improvement in the relative uncertainty on the integrated Compton parameter value $\Yv$. The latter highlights the utility of high resolution tSZ follow-up of \planck-discovered clusters to better constrain the $Y$--$M$ scaling relation used for cosmology studies based on cluster counts \citep{ProceedingMoriond}. We further combined the NIKA+\planck\ deprojected non-parametric pressure profile with the deprojected electronic density profile obtained from \xmm\ observations. This allowed us to obtain temperature and entropy profiles without recourse to X-ray spectroscopy and to undertake an hydrostatic mass analysis. The X-ray only (including spectroscopy) and the tSZ$+$X-ray (without spectroscopy) constraints are consistent within their uncertainties. This shows that high resolution tSZ observations, combined with X-ray snapshot imagery, are a competitive alternative to constrain cluster thermodynamics at high redshift, where X-ray spectroscopy requires large integration times to derive accurate temperature estimates. Comparison of the thermodynamic profiles to those obtained from the representative X-ray sample \rexcess\ \citep{boe07,universal,entropy_REXCESS}, in particular the radial distributions of temperature and entropy, indicates that \psz\ is a cool-core cluster. This result illustrates the complementarity between tSZ and X-ray data when only X-ray imaging observations are available.\\ The NIKA2 camera now installed at the focal plane of the IRAM 30-m telescope is currently undergoing commissioning. The number of detectors has been increased by a factor 10 with respect to the NIKA prototype to fully sample the telescope field of view of 6.5~arcmin. The NIKA2 tSZ Guaranteed Time Large Program \citep{NIKA2LP} is a follow-up of 50 SZ-discovered clusters with redshift up to $z = 1$, selected from the \planck\ and ACT catalogs \citep{ACT_cluster,Planck_cata2}. Following the work presented in this paper and in the previous NIKA studies \citep{RXJ1347NIKA,CLJ1227NIKA,MACSJ1424NIKA}, NIKA2 is expected to provide reliable tSZ detection and mapping of galaxy clusters in only a few hours integration time per cluster. Although NIKA2 alone will be a key tool for further understanding cluster physics, using the complementarity between different observational probes constitutes the best road for getting a comprehensive picture of the ICM. The NIKA2 data will therefore be complemented with ancillary data including X-ray, optical, and radio observations. The full data set will lead to significant improvements on the use of galaxy clusters to obtain constraints on cosmology and on the matter distribution and content of the Universe.
16
7
1607.07679
The determination of the thermodynamic properties of clusters of galaxies at intermediate and high redshift can bring new insights into the formation of large-scale structures. It is essential for a robust calibration of the mass-observable scaling relations and their scatter, which are key ingredients for precise cosmology using cluster statistics. Here we illustrate an application of high resolution (&lt;20 arcsec) thermal Sunyaev-Zel'dovich (tSZ) observations by probing the intracluster medium (ICM) of the Planck-discovered galaxy cluster PSZ1 G045.85+57.71 at redshift z = 0.61, using tSZ data obtained with the NIKA camera, which is a dual-band (150 and 260 GHz) instrument operated at the IRAM 30-m telescope. We deproject jointly NIKA and Planck data to extract the electronic pressure distribution from the cluster core (R 0.02 R<SUB>500</SUB>) to its outskirts (R 3 R<SUB>500</SUB>) non-parametrically for the first time at intermediate redshift. The constraints on the resulting pressure profile allow us to reduce the relative uncertainty on the integrated Compton parameter by a factor of two compared to the Planck value. Combining the tSZ data and the deprojected electronic density profile from XMM-Newton allows us to undertake a hydrostatic mass analysis, for which we study the impact of a spherical model assumption on the total mass estimate. We also investigate the radial temperature and entropy distributions. These data indicate that PSZ1 G045.85+57.71 is a massive (M<SUB>500</SUB> 5.5 × 10<SUP>14</SUP>M<SUB>⊙</SUB>) cool-core cluster. This work is part of a pilot study aiming at optimizing the treatment of the NIKA2 tSZ large program dedicated to the follow-up of SZ-discovered clusters at intermediate and high redshifts. This study illustrates the potential of NIKA2 to put constraints on thethermodynamic properties and tSZ-scaling relations of these clusters, and demonstrates the excellent synergy between tSZ and X-ray observations of similar angular resolution.
false
[ "intermediate redshift", "cluster statistics", "clusters", "R", "similar angular resolution", "X-ray observations", "new insights", "large program", "intermediate and high redshift", "intermediate and high redshifts", "PSZ1 G045.85", "SZ-discovered clusters", "R 0.02 R", "data", "the cluster core", "precise cosmology", "galaxies", "entropy distributions", "the Planck-discovered galaxy cluster", "key ingredients" ]
10.630835
3.365271
-1
12509796
[ "Singh, Avneet", "Papa, Maria Alessandra", "Eggenstein, Heinz-Bernd", "Zhu, Sylvia", "Pletsch, Holger", "Allen, Bruce", "Bock, Oliver", "Maschenchalk, Bernd", "Prix, Reinhard", "Siemens, Xavier" ]
2016PhRvD..94f4061S
[ "Results of an all-sky high-frequency Einstein@Home search for continuous gravitational waves in LIGO's fifth science run" ]
13
[ "Max-Planck-Institut für Gravitationsphysik, am Mühlenberg 1, 14476 Potsdam-Golm, Germany; Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; Leibniz Universität Hannover, Welfengarten 1, 30167 Hannover, Germany", "Max-Planck-Institut für Gravitationsphysik, am Mühlenberg 1, 14476 Potsdam-Golm, Germany; Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; University of Wisconsin-Milwaukee, Milwaukee, Wisconsin 53201, USA", "Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; Leibniz Universität Hannover, Welfengarten 1, 30167 Hannover, Germany", "Max-Planck-Institut für Gravitationsphysik, am Mühlenberg 1, 14476 Potsdam-Golm, Germany; Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany", "Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; Leibniz Universität Hannover, Welfengarten 1, 30167 Hannover, Germany", "Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; University of Wisconsin-Milwaukee, Milwaukee, Wisconsin 53201, USA; Leibniz Universität Hannover, Welfengarten 1, 30167 Hannover, Germany", "Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; Leibniz Universität Hannover, Welfengarten 1, 30167 Hannover, Germany", "Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; Leibniz Universität Hannover, Welfengarten 1, 30167 Hannover, Germany", "Max-Planck-Institut für Gravitationsphysik, Callinstraße 38, 30167 Hannover, Germany; Leibniz Universität Hannover, Welfengarten 1, 30167 Hannover, Germany", "University of Wisconsin-Milwaukee, Milwaukee, Wisconsin 53201, USA" ]
[ "2016PhRvD..94h2008Z", "2016PhRvD..94j2002A", "2016PhRvD..94l2006P", "2017PhRvD..96h2003S", "2017PhRvD..96l2004A", "2017PhRvD..96l4007Z", "2018PhRvD..97j2003A", "2018PhRvD..98h4058D", "2019PhRvD..99d4006S", "2021ApJ...909...79S", "2023APh...15302880W", "2023PhRvD.107f2005M", "2023PhRvD.108l2001D" ]
[ "astronomy", "physics" ]
12
[ "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "2005PhRvD..72f3006C", "2007PhRvD..76h2001A", "2008PhRvD..77b2001A", "2008PhRvD..78j2005P", "2009PhRvL.102k1102A", "2009RPPh...72g6901A", "2010NIMPA.624..223A", "2010PhRvD..82d2002P", "2011PhRvD..84b3007P", "2012PhRvD..85b2001A", "2012PhRvD..85d2003W", "2013PhRvD..87d2001A", "2013PhRvD..88d4004J", "2013PhRvD..88j2002A", "2016PhRvD..93f4011M", "2016PhRvD..93h4024K", "2016PhRvD..94d2002A", "2016PhRvD..94j2002A", "2017PhRvD..95b4022S" ]
[ "10.1103/PhysRevD.94.064061", "10.48550/arXiv.1607.00745" ]
1607
1607.00745_arXiv.txt
\label{section:intro} Ground-based gravitational wave (GW) detectors will be able to detect a continuous gravitational wave signal from a spinning deformed compact object provided that it is spinning with a rotational period between roughly $1$ and $100$ milliseconds, that it is sufficiently close to Earth and it is sufficiently ``bumpy''. Blind searches for continuous gravitational waves probe the whole sky and broad frequency ranges, looking for this type of objects. In this paper, we present the results of an all-sky \EatH search for continuous, nearly monochromatic, high-frequency gravitational waves in data from LIGO's 5th Science Run ($\mathsf{S5}$). A number of searches have been carried out on LIGO data \citep{S6BucketStage0,EarlyS5Paper,FullS5Semicoherent,S5EHHough,S4IncoherentPaper,S2FstatPaper} targeting lower frequency ranges. The only other search covering frequencies up to 1500$\,$Hz was conducted on $\mathsf{S6}$ data \citep{S6PowerFlux} taken at least 3 years apart from the data used here. Our search results are only 33\% less sensitive than those of \citet{S6PowerFlux}, even though the $\mathsf{S5}$ data is less sensitive than the $\mathsf{S6}$ data by more than a factor of 2. The search method presented here anticipates the procedure that will be used on the advanced detector (aLIGO) data. This search can be considered an extension of the $\mathsf{S5}$ \EatH search \cite{S5EHHough} although it employs a different search technique: this search uses the {\fontfamily{ppl}\selectfont\textit{Global Correlation Transform}} (GCT) method to combine results from coherent $\mathcal{F}${\fontfamily{ppl}\selectfont\textit{-statistic}} searches \citep{Holger2008,Holger2010}, as opposed to the previous \EatH search \cite{S5EHHough} that employed the {\fontfamily{ppl}\selectfont\textit{Hough-transform}} method to perform this combination. In the end, at fixed computing resources, these two search methods are comparable in sensitivity. However, a semi-coherent $\mathcal{F}${\fontfamily{ppl}\selectfont\textit{-statistic}} search is more efficient when considering a broad spin-down range, and for the \EatH searches we have decided to adopt it as our ``work horse''. We do not find any significant signal(s) among the set of searched waveforms. Thus, we set 90\%-confidence upper-limits on continuous gravitational wave strain amplitudes; near the lower end of the search frequency range between 1253.217--1255.217$\,$Hz, the most constraining upper-limit is $5.0\times 10^{-24}$, while toward the higher end of the search frequency range nearing 1500$\,$Hz, the upper-limit value is roughly $6.2\times 10^{-24}$. Based on these upper-limits, we can exclude certain combinations of signal frequency, star deformation (ellipticity) and distance values. We show with this search that even with $\mathsf{S5}$ data from the first generation of GW detectors, such constraints do probe interesting regions of source parameter space.
16
7
1607.00745
We present results of a high-frequency all-sky search for continuous gravitational waves from isolated compact objects in LIGO's fifth science run (S5) data, using the computing power of the Einstein@Home volunteer computing project. This is the only dedicated continuous gravitational wave search that probes this high-frequency range on S5 data. We find no significant candidate signal, so we set 90% confidence level upper limits on continuous gravitational wave strain amplitudes. At the lower end of the search frequency range, around 1250 Hz, the most constraining upper limit is 5.0 ×10<SUP>-24</SUP>, while at the higher end, around 1500 Hz, it is 6.2 ×10<SUP>-24</SUP>. Based on these upper limits, and assuming a fiducial value of the principal moment of inertia of 10<SUP>38</SUP> kg m<SUP>2</SUP> , we can exclude objects with ellipticities higher than roughly 2.8 ×10<SUP>-7</SUP> within 100 pc of Earth with rotation periods between 1.3 and 1.6 milliseconds.
false
[ "continuous gravitational wave strain amplitudes", "continuous gravitational waves", "the Einstein@Home volunteer computing project", "S5 data", "isolated compact objects", "rotation periods", "objects", "the only dedicated continuous gravitational wave search", "90% confidence level upper limits", "the computing power", "Hz", "ellipticities", "S5", "the search frequency range", "Earth", "inertia", "fifth", "LIGO", "these upper limits", "the higher end" ]
7.157747
2.261225
62
12517024
[ "Leahy, D.", "Ranasinghe, S." ]
2016sros.confE..86L
[ "HI absorption spectra for supernova remnants in the VGPS survey" ]
0
[ "University of Calgary, Canada;", "-" ]
null
[ "astronomy" ]
2
[ "Astrophysics - Astrophysics of Galaxies" ]
null
[ "10.48550/arXiv.1607.00059" ]
1607
1607.00059_arXiv.txt
$\,\!$\indent Supernova remnants (SNR) are an important research topic for astrophysics. They yield valuable information relevant to stellar evolution, the evolution of the Galaxy, and its interstellar medium. SNRs are the dominant source of kinetic energy input into the interstellar medium. A SNR emits in X-rays from its hot shocked gas, with temperature of order 1 keV. It emits radio continuum emission from relativistic electrons accelerated at the SNR shock. Finding the distance to a SNR is needed to determine its physical properties, such as luminosity, size, and age. The distance to a SNR can be estimated using its H I absorption spectrum combined with a model of Galactic rotation. In this paper we report the initial results from a study of the supernova remnants in the VLA Galactic Plane Survey (VGPS, Stil et al, 2006).
$\,\!$\indent Fig. 1 shows HI spectra for three different source/background areas for the SNR G41.1-0.3. The top (red) curves are the spectra from the background region. The middle (black) curves are the spectra from the source region, and lowest curves in each panel are the difference spectra from the source region. The spectra were extracted from the HI data cubes using task 'meanlev' from the DRAO Export Package Software. 'meanlev' allows the user to specify source and background regions in the HI data cube by using the radio continuum brightness level in the 1420MHz continuum image. This results in better separation of source and background regions than other methods which rely on specifying a geometrical shape in the HI data cube. The net result is a better HI absorption spectrum for any given data set. The current study has yielded a number of kinematic velocities for SNR in the VGPS survey area. For other SNR, we have obtained upper and/or lower limits to kinematic velocities. The velocities can be converted to distances using a rotation curve model for the Galaxy. We use the recent rotation curve published by Reid et al (2014). This is based mainly on very long baseline interferometry measurements of distances to a number of objects in the Galaxy. The full details of analysis for each SNR in the VGPS area, and comparison with previously published results, is currently being prepared for publication and release. This includes new (revised or improved) distances to many supernova remnants. \begin{figure} \center \includegraphics[width=\textwidth]{region_1_2.pdf} \caption{Sample HI spectra for the 3 different source/background regions for the supernova remnant G41.1-0.3 (also known as 3C 397). See text for details.} \end{figure} \small %
16
7
1607.00059
The set of supernova remnants (SNR) from Green's SNR catalog which are found in the VLA Galactic Plane Survey are studied here. For this set of 34 SNR, we extract and analyze HI absorption spectra in a uniform way and construct a catalog of absorption spectra and distance determinations. The results of this work will be summarized.
false
[ "HI absorption spectra", "absorption spectra and distance determinations", "the VLA Galactic Plane Survey", "SNR", "Greens SNR catalog", "supernova remnants", "Green", "a uniform way", "a catalog", "34 SNR", "this work", "The results", "The set", "this set", "34", "we", "which" ]
4.113352
5.026468
60
2323953
[ "Garraffo, Cecilia", "Drake, Jeremy J.", "Cohen, Ofer" ]
2016A&A...595A.110G
[ "The missing magnetic morphology term in stellar rotation evolution" ]
66
[ "Harvard-Smithsonian Center for Astrophysics, 60 Garden St. Cambridge, MA, 02138, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden St. Cambridge, MA, 02138, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden St. Cambridge, MA, 02138, USA" ]
[ "2016ApJ...832...97S", "2017ApJ...845...46F", "2017ApJ...849...83P", "2017ApJ...850...45R", "2017IAUS..328..315A", "2017MNRAS.472..876O", "2017SoPh..292..126M", "2018ApJ...854...78F", "2018ApJ...856...53P", "2018ApJ...862...90G", "2018ApJ...862...93A", "2018ApJ...864..125F", "2018ApJ...868...60G", "2019A&A...626A..22J", "2019ApJ...871...39M", "2019ApJ...876...44F", "2019ApJ...879...12S", "2019ApJ...883...67F", "2019ApJ...886..120S", "2019ApJ...887..199G", "2019LNP...955.....L", "2019MNRAS.483.1125J", "2019MNRAS.485.3661F", "2020A&A...635A.170A", "2020A&A...641A..51F", "2020A&A...643A.129P", "2020ApJ...893..107B", "2020ApJ...894...69S", "2020ApJ...895...47A", "2020ApJ...895..140H", "2020ApJ...905..107M", "2020JGRA..12527639G", "2020MNRAS.491..455B", "2020arXiv200303231G", "2021A&A...649A..96J", "2021ApJ...912...65G", "2021ApJ...916...66B", "2021ApJ...921..122M", "2021LRSP...18....3V", "2021MNRAS.500.3438O", "2021MNRAS.500.4560P", "2021MNRAS.506.2309E", "2021NatAs...5..707H", "2022ApJ...925..100I", "2022ApJ...933..114D", "2022ApJ...933L..17M", "2022ApJ...941L...8G", "2022MNRAS.517.4916E", "2022csss.confE.199G", "2023ApJ...949...54C", "2023ApJ...950...27G", "2023ApJ...951...79B", "2023ApJ...957...71S", "2023ApJ...959..140M", "2023MNRAS.518.1683K", "2023MNRAS.525.2708C", "2023MNRAS.526..870S", "2023SSRv..219...70I", "2024A&A...681A..83G", "2024ApJ...962...54W", "2024ApJ...962..138S", "2024FrASS..1156379S", "2024SSRv..220....7G", "2024arXiv240205912S", "2024arXiv240317279S", "2024arXiv240500779S" ]
[ "astronomy" ]
10
[ "magnetic fields", "magnetohydrodynamics (MHD)", "stars: activity", "stars: rotation", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1962AnAp...25...18S", "1967ApJ...148..217W", "1967ApJ...150..551K", "1968MNRAS.138..359M", "1969SoPh....9..131A", "1981ApJ...243..625E", "1984A&A...136...98M", "1984ApJ...280..202S", "1987ApJ...318..337S", "1988ApJ...333..236K", "1993ApJ...409..624S", "1995GeoRL..22.3301P", "1998A&A...335..183Q", "1999JCoPh.154..284P", "2000AJ....119.1303T", "2003ApJ...586..464B", "2005ApJ...634L.193A", "2007MNRAS.377.1488H", "2007RvGeo..45.1004M", "2009ARA&A..47..333D", "2009ApJ...707.1372W", "2009MNRAS.398.1383F", "2011ApJ...733..115M", "2011MNRAS.413.1922M", "2011MNRAS.413.1949W", "2011MNRAS.417.2592C", "2012ApJ...754L..26M", "2012JCoPh.231..870T", "2013ApJ...764...23S", "2013ApJ...764...32G", "2013ApJ...778..176O", "2013MNRAS.435.1451F", "2014ASTRP...1...43L", "2014ApJ...782...81V", "2014ApJ...783...55C", "2014ApJ...789..101B", "2014JGRA..119...11J", "2014MNRAS.438.1162V", "2015ApJ...798..116R", "2015ApJ...813...40G", "2015ApJ...814...99R", "2015MNRAS.449....8W", "2016MNRAS.457..580F" ]
[ "10.1051/0004-6361/201628367", "10.48550/arXiv.1607.06096" ]
1607
1607.06096_arXiv.txt
\label{sec:Intro} One of the fundamental observable characteristics of a star is its rotation period. Stellar rotation evolves over time as a result of interior structural adjustments as stars settle onto the main sequence, and as a result of angular momentum loss through magnetized stellar winds \citep[e.g.][]{Schatzman:62,Kraft:67,Weber.Davis:67,Mestel:68,Endal.Sofia:81,Kawaler:88}. However, our present understanding of the details of this behavior is far from complete. Observations of young open clusters \cite[e.g.][see \citealt{Meibom.etal:11} for a recent compilation]{Stauffer.Hartmann:87, Soderblom.etal:93, Queloz.etal:98, Terndrup.etal:00} have revealed a bimodal distribution of fast and slower rotation rates that has proven difficult to explain with current spin-down models \cite[e.g.][]{Stauffer.etal:84,Soderblom.etal:93, Barnes:03}. The situation has been summarized recently by \citet{Brown:14}. Recently, the morphology of the magnetic field (by morphology we mean the distribution of the magnetic field on the stellar surface, which some authors call topology) has received a lot of attention in the context of stellar rotation evolution \citep{Brown:14, Reville.etal:15a, Garraffo.etal:15, Reville.etal:15b}. Most previous spin down models had assumed dipolar morphology, with few exceptions that explored the role of the multipole order of the magnetic field \citep{Weber.Davis:67, Mestel.Paris:84, Kawaler:88} but were limited to the effect of the radial dependence of the magnetic field. A growing number of studies based on ZDI maps indicate that young, active stars have a larger fraction of their magnetic energy stored in higher order multipole components \citep[e.g.][]{Donati:03, Donati.Landstreet:09, Marsden.etal:11a, Waite.etal:11, Waite.etal:15, Folsom.etal:16}. \cite{Linsky.Wood:14} have also recently inferred mass loss rates for the active stars $\xi$~Boo~A and $\pi^1$~UMa that are two orders of magnitude lower than expected based on extrapolation from lower activity stars, suggesting that magnetic topology could have a more profound effect on angular momentum loss than simply through the radial field strength dependence. \citet[][from here on CG15]{Garraffo.etal:15} used detailed three-dimensional magnetohydrodynamic (MHD) stellar wind simulations to show that a factor representing magnetic morphology, missing in rotation evolution models, can have a drastic effect on mass and angular momentum losses. They drew a connection between magnetic morphology and the ``metastable dynamo" proposal of \citet{Brown:14} in which bimodal rotation distributions at early ages are attributed to an initially weak coupling of magnetic field to the stellar wind that ``spontaneously and randomly'' changes to a strongly coupled mode and initiation of rapid spin-down. CG15 suggested that magnetic morphology could provide such a mode switch. That work was based on a limited set of cases of magnetic complexity assuming idealized magnetic field distributions. Here, we investigate the effects of different flux distributions (different $m$) of each term $n$ in a spherical harmonic multipolar expansion representation of the surface magnetic field. Based on these complete study of magnetic flux complexity and distribution, we derive the analytical dependence of mass and angular momentum loss rates on magnetic complexity, which provides the means to realistically estimate those rates without the need of computationally expensive simulations. The numerical methods are described in Section~\ref{sec:Methods} and the results of model calculations in Section~\ref{sec:Results}. We state our main findings in Section~\ref{sec:Conclusions} .
\label{sec:Conclusions} Our MHD simulations indicate that mass and angular momentum loss rates from solar-like stars are strongly suppressed by magnetic complexity but independently of how this complexity is organized. The loss rates are controlled by the number of rings in which the polarity changes sign, labeled by the spherical harmonic order of complexity, $n$. Translating these results to non-ideal surface field distributions of real stars, both mass and angular momentum loss rates depend on the level of complexity of the field only and not on the details of the field distribution over the stellar surface. We have provided analytical formulae for the dependence of mass and angular momentum loss rates on magnetic complexity that can be applied in stellar rotation evolution models. We have shown that this ingredient significantly improves the analytical estimations of mass and angular momentum loss rates based on ZDI maps (total magnetic flux and complexity) over the usual dipolar assumption.
16
7
1607.06096
<BR /> Aims: This study examines the relationship between magnetic field complexity and mass and angular momentum losses. Observations of open clusters have revealed a bimodal distribution of the rotation periods of solar-like stars that has proven difficult to explain under the existing rubric of magnetic braking. Recent studies suggest that magnetic complexity can play an important role in controlling stellar spin-down rates. However, magnetic morphology is still neglected in most rotation evolution models due to the difficulty of properly accounting for its effects on wind driving and angular momentum loss. <BR /> Methods: Using state-of-the-art magnetohydrodynamical magnetized wind simulations we study the effect that different distributions of the magnetic flux at different levels of geometrical complexity have on mass and angular momentum loss rates. <BR /> Results: Angular momentum loss rates depend strongly on the level of complexity of the field but are independent of the way this complexity is distributed. We deduce the analytical terms representing the magnetic field morphology dependence of mass and angular momentum loss rates. We also define a parameter that best represents complexity for real stars. As a test, we use these analytical methods to estimate mass and angular momentum loss rates for 8 stars with observed magnetograms and compare them to the simulated results. <BR /> Conclusions: Magnetic field complexity provides a natural physical basis for stellar rotation evolution models requiring a rapid transition between weak and strong spin-down modes.
false
[ "Angular momentum loss rates", "angular momentum loss", "Magnetic field complexity", "magnetic field complexity", "angular momentum", "magnetic complexity", "mass and angular momentum loss rates", "geometrical complexity", "losses", "complexity", "stellar rotation evolution models", "observed magnetograms", "magnetic braking", "real stars", "different levels", "most rotation evolution models", "magnetic morphology", "different distributions", "mass", "wind driving" ]
8.661264
12.052974
87
12594325
[ "Rácz, Gábor", "Dobos, László", "Beck, Róbert", "Szapudi, István", "Csabai, István" ]
2017MNRAS.469L...1R
[ "Concordance cosmology without dark energy" ]
48
[ "Department of Physics of Complex Systems, Eötvös Loránd University, Pf. 32, H-1518 Budapest, Hungary", "Department of Physics of Complex Systems, Eötvös Loránd University, Pf. 32, H-1518 Budapest, Hungary", "Department of Physics of Complex Systems, Eötvös Loránd University, Pf. 32, H-1518 Budapest, Hungary", "Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu, HI 96822, USA", "Department of Physics of Complex Systems, Eötvös Loránd University, Pf. 32, H-1518 Budapest, Hungary" ]
[ "2016arXiv161207245C", "2017A&A...598A.111R", "2017ApJ...841...63T", "2017CQGra..34w7001D", "2017JCAP...02..047S", "2017JCAP...06..025B", "2017JCAP...12..008M", "2017MNRAS.469..744K", "2017MNRAS.470.4493W", "2017PhRvD..95l4009F", "2018A&A...610A..51R", "2018CQGra..35b4003B", "2018MNRAS.473L..46B", "2018MNRAS.475.1777K", "2018MNRAS.476..451O", "2018MNRAS.479.3582B", "2018arXiv180105033O", "2019A&A...622A..83S", "2019A&A...624A..12W", "2019AN....340..586K", "2019CQGra..36d5003B", "2019JCAP...10..036K", "2019JCAP...12..049R", "2019MNRAS.484.5267K", "2019PhRvD..99h3521S", "2020AdAst2020E..17Y", "2020Ap&SS.365...16C", "2020ApJ...897..166L", "2020MNRAS.491.2330F", "2020MNRAS.493.4830A", "2020MNRAS.499..320K", "2020PhyS...95f5005D", "2021EPJC...81..208S", "2021MNRAS.500.3838D", "2021MNRAS.501.1481H", "2022CQGra..39u5007B", "2022CQGra..39v5008B", "2022ChJPh..77.1732S", "2022Entrp..24..976S", "2022JHEAp..34...49A", "2022MNRAS.513...15K", "2022NewAR..9501659P", "2023ApJ...959...83L", "2023PhRvD.107l3536M", "2023Univ....9..302H", "2023arXiv231019451V", "2023arXiv231101438L", "2024Univ...10..119K" ]
[ "astronomy" ]
12
[ "methods: numerical", "cosmological parameters", "dark energy", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnrasl/slx026", "10.48550/arXiv.1607.08797" ]
1607
1607.08797_arXiv.txt
Gravitation being the only effective force on the largest scales, cosmological evolution is governed by general relativity (GR). To zeroth order, the homogeneous and isotropic Friedmann--Lema\^{\i}tre--Robertson--Walker (FLRW) solutions of Einstein's equations drive the expansion and growth history of the Universe. The concordance $\Lambda$CDM model \citep[e.g.,][]{Planck2015} posits an unknown form of energy with negative pressure and an energy density about $10^{123}$ off from theoretical expectations. The $\Lambda$CDM paradigm reproduces most observations, although, to this day, no plausible candidate for dark energy has emerged and some tensions remain \citep[see][for a recent comprehensive review]{Buchert_2016IJMPD..2530007B}. Most notably, the latest local measurements \citep{Riess_hubbleParameterTension_2016} of the Hubble constant are up to $3.4\sigma$ high compared to the value derived from \textit{Planck} observations \citep{Planck2015} of the cosmic microwave background. The ubiquitous presence of clusters, filaments, and voids in the cosmic web manifestly violate the assumed homogeneity of $\Lambda$CDM. Given the non-linear nature of Einstein's equations, it has been known for a while that local inhomogeneities influence the overall expansion rate, whereas the magnitude of such backreaction effect is debated. In particular, \citet{2014CQGra..31w4003G,2016CQGra..33l5027G} argued that the effect of inhomogeneities on the expansion of the Universe is irrelevant, while \citet{2015CQGra..32u5021B} disputed the general applicability of the former proof. More recently, \citet{PhysRevLett.116.251301} used numerical relativity to show the existence of a departure from FLRW behaviour due to inhomogeneities, beyond what is expected from linear perturbation theory. Nevertheless, the spectacular successes of the homogeneous concordance model suggest that any effect of the inhomogeneities on the expansion rate should be weak, {\em unless} it mimics the $\Lambda$CDM expansion and growth history to a degree allowed by state of the art observations. In this spirit, we present a statistical non-perturbative algorithm, a simple modification to standard $N$-body simulations, that provides a viable alternative to dark energy while it can simultaneously resolve the Hubble constant puzzle. In the late-time non-linear evolution of the Universe, coarse graining and averaging are both problematic \citep[see][and references therein]{Wiltshire2007a, WiltshireCQG}. The complexity of Einstein's equations prevents direct numerical modelling of backreaction, which ideally would require a general relativistic simulation of space-time seeded with small initial fluctuations. Such an ideal simulation would contain a hierarchy of coarse graining scales describing the space-time of stars, galaxies, galaxy clusters, intra-cluster medium and dark matter, etc., and their metric would be stitched together in a careful manner \citep[see][for a discussion of the hierarchy of coarse graining scales]{Wiltshire2014}. As fluctuations grow due to non-linear gravitational amplification, space-time itself becomes complex, and even the concept of averaging becomes non-trivial. Despite the difficulties, \citet{Buchert2000, Buchert2001} and others \citep{2012ARNPS..62...57B} realized that backreaction can be understood in a statistical fashion through the spatial averaging of Einstein's equations on a hypersurface, leading to the Buchert equations. Second-order perturbative solutions to the equations are given by \cite{Kolb2005} and \cite{Rasanen2010}, while \cite{Wiltshire2007a, Wiltshire2007b, Wiltshire2009} presents an analytical solution to a two-scale (voids and walls) inhomogeneous universe. We propose a non-perturbative, multi-scale statistical approach to study strong backreaction based on the general relativistic separate universe conjecture, which states that a spherically symmetric region in an isotropic universe behaves like a mini-universe with its own energy density $\Omega = 1 + \delta$ \citep{Weinberg2008}. The conjecture was proven by \cite{Dai2015} for compensated top hat over- and underdensities. In the quasi Newtonian framework, the separate universe conjecture is widely used in successful spherical collapse models \citep{Bernardeau1994, Mohayaee2006, 2016MNRAS.455L..11N}. We build on this conjecture to estimate the expansion rate as the volume average of local expansion rates, avoiding the calculation of any geometric quantities such as the average curvature of the universe. Our algorithm neglects tidal forces, but the scales over which averages are calculated are solidly grounded in Newtonian physics, simplifying the interpretation of our results. We show that under our algorithm, the non-Gaussian distribution of matter arising from the non-linearities of the cosmological fluid equations causes the expansion rate to decrease at a slower rate than normally calculated from the global Friedmann equations, thereby mimicking the effect of dark energy. In our scheme the coarse graining scale is an adjustable, phenomenological parameter corresponding to the best ``particle size'' to use when modelling the evolution of the Universe. When the coarse graining scale approaches the scale of homogeneity, our model, obviously, shows no effect: in this limit it is equivalent to the global Friedmann equations. Approaching very small scales, the assumptions of the model progressively break down due to the increasing anisotropy around, and inhomogeneity inside, the spherically symmetric regions. Somewhere between the extremes, there is an optimal scale that we expect to be around the size of virialized structures detached from the Hubble flow, therefore on the order of $10^9-10^{13} M_{\odot}$. The coarse graining scale is a semi-nuisance parameter to be fit, analogous to halo model parameters. While in this work we use a single, redshift-independent, comoving coarse graining scale, in principle, the optimal scale could depend on the state of the Universe and its constituents and thus, on redshift.
We have presented a modified $N$-body simulation where we estimated the global expansion rate by averaging local expansion rates of mini-universes based on the separate universe approximation. While we do not attempt to connect space-time regions or compute light propagation across curved regions, our approach is equivalent to a non-perturbative statistical backreaction calculation. Our approximation neglects tidal forces due to anisotropies, and has an ambiguity associated with the optimal coarse graining scale. For a large enough scale, the effect is negligible, while on small scales anisotropies break the underlying assumptions of the approximation. Since virialised objects detach from the expansion, we expect that the optimal coarse graining scale, treated as a nuisance parameter, is related to the size of the typical virialised regions. \begin{table} \begin{tabular}{r | c | c} $N$ & $M \left[ 10^{11} M_\odot \right]$ & $H_0~\left[ \Hunit \right]$ \\ \hline 135,000 & 9.40 & 65.4 \\ 320,000 & 3.96 & 68.9 \\ 625,000 & 2.03 & 71.4 \\ 1,080,000 & 1.17 & 73.1 \\ \end{tabular} \caption{Summary of simulation input parameters and the resulting values of $H_0$. In all cases the linear size of the simulation box was $L=147.62~\unit{Mpc}$ and the early epoch value of the Hubble parameter was set to $H_{z=9} = 1191.9~\Hunit$, complying with \textit{Planck} $\Lambda$CDM best-fit parameters.} \label{tab:parameters} \end{table} Our modified $\Omega_m=1$ simulation mimics the $\Lambda$CDM expansion history remarkably well. Since growth history is also driven by the expansion history, we expect that our simulations are consistent with luminosity distance and Hubble parameter observations constraining dark energy. Present-day supernova observations are well fit by our model, moreover, our model naturally resolves the tension between local and CMB Hubble constant measurements. Detailed fits to observations, and forecasting for future surveys such as Euclid, WFIRST, HSC, etc. is left for future work, but it is clear already from Fig.~\ref{fig.at_1200} that if our model is sufficiently different from the standard $w=-1$ vacuum energy model, upcoming surveys will be able to confirm or rule it out. We also note that some of our results are numerically very similar to the analytically derived timescape scenario presented in \cite{Wiltshire2007b, Wiltshire2009}, in addition to sharing the separate universe approximation. Further investigation is yet to be done to compare the two approaches in detail. To investigate the validity of the AvERA simulation, we performed follow-up analytic calculations. If inhomogeneities mainly affect structure growth via the expansion history of the universe, the (near) lognormal approximation found in $\Lambda$CDM simulations should approximate well the density distribution of the mini-universes. For instance, to calculate the longitudinal comoving scale, we calculate the average \begin{equation} \frac{1}{H_\textnormal{eff}(z)} = \avg{\frac{1}{H(z)}} = \int_{-1}^{\infty} P(1+\delta) \frac{1}{H_{\mathcal{D}}(1+\delta, z)} \,\textnormal{d}\delta, \end{equation} where $P(1+\delta)$ is a lognormal distribution, and $H_{\mathcal{D}}(1+\delta, z)$ is calculated from Equation~\ref{eq:local_Friedamnn_eq} with $\Omega = 1+\delta$. The variance of the lognormal distribution is estimated from $\sigma_A^2 = 0.73\log (1+\sigma^2_{lin}/0.73)$ \citep{ReppSzapudi2017}, and $\avg{\log(1+\delta)} = - 0.67 \log(1+\sigma^2_\textnormal{lin}/(2\times 0.67))$ \citep{ReppSzapudi2017}, where $\sigma_\textnormal{lin}$ is the linear variance of dark matter fluctuations on the coarse graining scale. The lognormal PDF is a good approximation and the above fits are accurate for concordance cosmologies. The result of such a calculation is shown on the right panel of Figure~\ref{fig:OmegaEvol}. Details will be presented elsewhere (Szapudi et.al. 2017 in prep.). The theoretical calculation is insensitive to the details of the PDF, i.e. departures from lognormality. We can calculate the effective expansion rate fairly accurately by replacing $\Omega_\textnormal{M}$ with its most likely value, i.e. $1+\delta$ at the peak of the density PDF on the left panel of Figure~\ref{fig:OmegaEvol}. This corresponds to approximating the lognormal PDF with a Dirac-$\delta$ function centred on the peak of the distribution. Thus the physical meaning of these calculations is simple: according to our approximation, it is not the average but the typical energy density that governs the expansion rate of the Universe. At high redshifts, where the distribution is fairly symmetric, the typical value of $\delta$ (mode of the PDF) is close to the average and the Universe evolves without backreaction. At late times skewness increases, the volume of the Universe is dominated by voids, and the typical value of $\delta$ is negative, thus effectively $\Omega_\textnormal{M} < 1$. High density regions, where metric perturbations are perhaps the largest, are inconsequential to this effect: what matters is the non-Gaussianity of the density distribution, in particular, the large volume fraction of low density regions, as advertised earlier. \begin{figure} \begin{center} \includegraphics[width=\columnwidth]{./fig/theorFig} \vspace*{-0.5cm} \end{center} \caption{Left: Evolution of the distribution of $1+\delta$ during the AvERA N-body simulation. Right: The normalized line of sight comoving distance $D_c^*= \int_0^z H_0/(H(z^\prime))dz^\prime$ at each time step has been calculated from a matter-only FLRW model with $\Omega_m$ corresponding to the actual peak $1+\delta$ (blue). The curve deviates from the $\Omega_m=1$ model (green) and closely follows the $\Lambda$CDM model (red).} \label{fig:OmegaEvol} \end{figure} The statistical approach we use is spatial (volume) averaging. While averaging is ambiguous in curved space-times \citep{Wiltshire2014}, note that all astrophysical quantities, most notably the power spectrum that is used to calculate all cosmological parameters, are estimated through analogous statistical procedures. We neglect local anisotropies, and we assume that those effects average out over time. Nevertheless, one could generalize our code to include tidal forces using elliptical collapse equations for the mini-universes we consider. This refinement of our calculations would quantify to first order the effect from tidal forces, but is left for future work. Our simple model with a reasonable coarse graining scale yields results that are consistent with the observations, and the model also has a simple physical interpretation. Further studies are needed both on the theoretical front and on fitting cosmological parameters, nevertheless, our approach is not only a viable alternative to dark energy models, but appears to be flexible enough to resolve some tensions in a natural way.
16
7
1607.08797
According to the separate universe conjecture, spherically symmetric sub-regions in an isotropic universe behave like mini-universes with their own cosmological parameters. This is an excellent approximation in both Newtonian and general relativistic theories. We estimate local expansion rates for a large number of such regions, and use a scale parameter calculated from the volume-averaged increments of local scale parameters at each time step in an otherwise standard cosmological N-body simulation. The particle mass, corresponding to a coarse graining scale, is an adjustable parameter. This mean field approximation neglects tidal forces and boundary effects, but it is the first step towards a non-perturbative statistical estimation of the effect of non-linear evolution of structure on the expansion rate. Using our algorithm, a simulation with an initial Ω<SUB>m</SUB> = 1 Einstein-de Sitter setting closely tracks the expansion and structure growth history of the Λ cold dark matter (ΛCDM) cosmology. Due to small but characteristic differences, our model can be distinguished from the ΛCDM model by future precision observations. Moreover, our model can resolve the emerging tension between local Hubble constant measurements and the Planck best-fitting cosmology. Further improvements to the simulation are necessary to investigate light propagation and confirm full consistency with cosmic microwave background observations.
false
[ "local scale parameters", "cosmic microwave background observations", "local expansion rates", "future precision observations", "non-linear evolution", "such regions", "local Hubble constant measurements", "full consistency", "structure", "boundary effects", "light propagation", "regions", "universes", "their own cosmological parameters", "a non-perturbative statistical estimation", "a scale parameter", "-", "the Λ cold dark matter", "an adjustable parameter", "the expansion rate" ]
10.95246
0.982867
-1
12408779
[ "Núñez, Alejandro", "Agüeros, Marcel A." ]
2016ApJ...830...44N
[ "The X-Ray Luminosity Function of M37 and the Evolution of Coronal Activity in Low-mass Stars" ]
13
[ "Department of Astronomy, Columbia University, 550 West 120th Street, New York, NY 10027, USA", "Department of Astronomy, Columbia University, 550 West 120th Street, New York, NY 10027, USA" ]
[ "2017ApJ...834..176N", "2018A&A...615A...2C", "2018ApJ...862...33A", "2018ApJ...863...91F", "2018PhDT........75N", "2020ApJ...898..104E", "2021A&A...649A..96J", "2021A&A...653A..49B", "2021A&A...654A.134C", "2022MNRAS.512...41K", "2024ApJ...960...62E", "2024ApJ...961...26R", "2024arXiv240613757W" ]
[ "astronomy" ]
11
[ "open clusters and associations: individual: M37", "stars: activity", "stars: coronae", "stars: low-mass", "X-rays: individual: M37", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.3847/0004-637X/830/1/44", "10.48550/arXiv.1607.05789" ]
1607
1607.05789_arXiv.txt
In low-mass stars ($\lapprox$1~\Msun), X rays originate in a magnetically heated corona, and serve as a proxy for the strength of the magnetic dynamo\footnote{For a recent review of the connection between stellar activity and coronal heating, see \citet{Testa2015}.}. The X-ray luminosity, \LX, is correlated with age and rotation \citep[e.g.,][]{palla81, pizzo03}, and decreases as low-mass stars spin down because of the loss of angular momentum through magnetized winds. Calibrating the evolution of \LX\ is key to quantifying the interplay between stellar rotation and magnetic fields, and ultimately to uncovering the still-mysterious processes responsible for these fields. Observations indicate that \LX\ does not decay smoothly with age ($t$). Surveys of solar-type stars in the Pleiades ($t \approx 125$ Myr) and the Orion Nebula Cluster (ONC; 1$-$10 Myr) concluded that \LX\ falls off relatively slowly early in a star's life: \LX\ $\propto t^{-0.76}$ \citep{queloz98,preibisch2005a}. Because there are few accessible $t\ \gapprox\ 200$ Myr clusters, constraints are harder to come by for older stars, but from observations of five solar analogs, \citet{guedel1997} determined that \LX~$\propto t^{-1.5}$ for $t > 1$ Gyr, as did \citet{giardino2008} from their survey of the $\approx$1.5-Gyr-old cluster NGC 752. Core-envelope decoupling or a change in the magnetic field topology are the commonly invoked explanations for this sharp drop off in \LX\ \citep[e.g.,][]{kawaler1988,Krishnamurthi1997}, but it remains poorly understood. Interestingly, chromospheric activity, another proxy for magnetic field strength, appears to evolve differently \citep[e.g.,][]{Jackson2010, Douglas14} but may suffer a similarly steep decline for $t>1$ Gyr \citep{pace2004}. To determine the evolution of \LX, we need more $\gapprox$200-Myr-old open clusters with well-characterized cumulative X-ray luminosity functions (XLFs). In \citet[][hereafter Paper I]{Nunez2015}, we described our 440.5 ks {\it Chandra} observation of M37 (NGC 2099), a rich, $\approx$500-Myr-old cluster at a distance of 1490$\pm$120 pc \citep{hartman08b}. We combined the photometry compiled by \citet[hereafter HA08]{Hartman2008a} and distance from the cluster core to generate membership probabilities (\Pmem) for cluster stars. Our final catalog included 561 X-ray sources with optical counterparts, 278 of which had \Pmem~$\geq 0.2$. Here, we add to these detections \LX\ upper limit (UL) measurements for undetected members to determine the XLFs for M37's G, K, and M stars. We also compute bolometric luminosities (\Lbol), and use these to determine the \LL\ functions (LLFs) for these stars, thereby establishing M37 as the benchmark 500-Myr-old cluster for studies of the evolution of X-ray emission in low-mass stars. In Section 2, we describe our optical and X-ray data for M37, outline how we assign membership thresholds for inclusion in our analysis, and calculate \LX\ ULs for undetected sources. In Section 3, we construct the XLFs and LLFs and discuss the impact of our upper limits and of binaries on these functions. In Section 4, we first homogenize the \LX, \Lbol, and masses of stars in six other clusters ranging in age from 6 to 625 Myr. We then examine the evolution of the XLFs and LLFs for GKM stars over $\approx$600 Myr. We present our conclusions in Section 5.
16
7
1607.05789
We use a 440.5 ks Chandra observation of the ≈500 Myr old open cluster M37 to derive the X-ray luminosity functions of its ≤1.2 {M}<SUB>⊙ </SUB> stars. Combining detections of 162 M37 members with upper limits for 160 non-detections, we find that its G, K, and M stars have a similar median (0.5-7 keV) X-ray luminosity {L}<SUB>{{X</SUB>}}={10}<SUP>29.0</SUP> {erg} {{{s}}}<SUP>-1</SUP>, whereas the {L}<SUB>{{X</SUB>}}-to-bolometric-luminosity ratio ({L}<SUB>{{X</SUB>}}/{L}<SUB>{bol</SUB>}) indicates that M stars are more active than G and K stars by ≈ 1 order of magnitude at 500 Myr. To characterize the evolution of magnetic activity in low-mass stars over their first ≈ 600 {{Myr}}, we consolidate X-ray and optical data from the literature for stars in six other open clusters: from youngest to oldest they are, the Orion Nebula Cluster (ONC), NGC 2547, NGC 2516, the Pleiades, NGC 6475, and the Hyades. For these, we homogenize the conversion of instrumental count rates to {L}<SUB>{{X</SUB>}} by applying the same one-temperature emission model as for M37, and obtain masses using the same empirical mass-absolute magnitude relation (except for the ONC). We find that for G and K stars X-ray activity decreases ≈ 2 orders of magnitude over their first 600 Myr, and for M stars, ≈1.5. The decay rate of the median {L}<SUB>{{X</SUB>}} follows the relation {L}<SUB>{{X</SUB>}}\propto {t}<SUP>b</SUP>, where b=-0.61+/- 0.12 for G stars, -0.82 ± 0.16 for K stars, and -0.40 ± 0.17 for M stars. In {L}<SUB>{{X</SUB>}}/{L}<SUB>{bol</SUB>} space, the slopes are -0.68 ± 0.12, -0.81 ± 0.19, and -0.61 ± 0.12, respectively. These results suggest that for low-mass stars the age-activity relation steepens after ≈ 625 {{Myr}}, consistent with the faster decay in activity observed in solar analogs at t\gt 1 {{Gyr}}.
false
[ "K stars", "G stars", "M stars", "stars", "G and K stars X-ray activity", "NGC", "-0.68 ±", "±", "-", "Myr", "the X-ray luminosity functions", "X", "low-mass stars", "≈", "the ≈500 Myr old open cluster M37", "K", "magnetic activity", "activity", "G", "magnitude" ]
8.535779
12.130387
87
12510330
[ "Hohmann, Manuel", "Järv, Laur", "Kuusk, Piret", "Randla, Erik", "Vilson, Ott" ]
2016PhRvD..94l4015H
[ "Post-Newtonian parameter γ for multiscalar-tensor gravity with a general potential" ]
30
[ "Institute of Physics, University of Tartu, W. Ostwaldi 1, Tartu 50411, Estonia", "Institute of Physics, University of Tartu, W. Ostwaldi 1, Tartu 50411, Estonia", "Institute of Physics, University of Tartu, W. Ostwaldi 1, Tartu 50411, Estonia", "Institute of Physics, University of Tartu, W. Ostwaldi 1, Tartu 50411, Estonia", "Institute of Physics, University of Tartu, W. Ostwaldi 1, Tartu 50411, Estonia" ]
[ "2016arXiv160903159C", "2017PhRvD..95h4050P", "2017PhRvD..96j4026H", "2017PhRvL.118o1302J", "2017SCPMA..60j0411Z", "2017Univ....3...37J", "2018JCAP...06..032C", "2018NuPhB.927..219K", "2018PDU....22..101S", "2018PhRvD..97j4011H", "2018PhRvD..98f4004H", "2019EPJP..134..318L", "2019IJMPD..2850132L", "2019PhLB..795..129L", "2019arXiv191108313C", "2020PhLB..81135985L", "2020PhRvD.101b4005F", "2020PhRvD.101b4017E", "2021PhRvD.103d4030F", "2021PhRvD.104l4030R", "2021arXiv210512582S", "2022IJMPD..3150117L", "2023CQGra..40m5001B", "2023EPJP..138..387M", "2023FoPh...53....5L", "2024PDU....4401459L", "2024PhRvD.109d4045A", "2024PhRvD.109d4070H", "2024PhRvD.109h4073J", "2024arXiv240400094N" ]
[ "astronomy", "physics" ]
3
[ "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]
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[ "10.1103/PhysRevD.94.124015", "10.48550/arXiv.1607.02356" ]
1607
1607.02356_arXiv.txt
Multiscalar-tensor gravity (MSTG) generalizes the scalar-tensor gravity (STG) of a scalar field $\Phi$ nonminimally coupled to curvature $R$, to the case of multiple scalar fields $\Phi^\alpha$ \cite{DamourEF,Berkin:1993bt}. Nonminimal couplings are typically generated by quantum corrections and arise in the effective models of higher dimensional theories. Diverse versions of MSTG appear in fundamental physics and cosmology in various constructions and under different disguises. First, there are several phenomenological motivations to consider nonminimally coupled scalars. The Standard Model Higgs field is an SU(2) complex doublet, in the case it is endowed with a nonminimal coupling to curvature also the Goldstone modes may play a role in Higgs inflation \cite{Higgs inflation} and the subsequent dark energy era \cite{Higgs dark energy}. Otherwise a nonminimal Higgs may be paired with another nonminimal scalar (e.g.\ a dilaton) \cite{Higgs and scalar}, or the inflation and dark energy could be run by two nonminimally coupled scalars \cite{biscalar models}. More general MSTG inflation or dark energy models have $N$ fields with noncanonical kinetic terms and arbitrary potential \cite{DamourEF,MSTG inflation,Einstein frame metric,Kuusk:2014sna,meie2015} (also considered for stellar models \cite{Horbatsch:2015bua}), or are embedded into a supergravity setup \cite{Higgs sugra}. The most general multiscalar-tensor gravitational action with second order field equations includes derivative couplings and is a generalization of Horndeski's class of theories, so far worked out for the two fields case \cite{biscalar Horndeski}. Second, different proposed extensions and modifications of general relativity can be also cast into the form of MSTG by a change of variables. It is well-known that, if the gravitational lagrangian is nonlinear in curvature, $f(R)$, or more generally $f(\Phi^\alpha,R)$, the theory is dynamically equivalent to (M)STG with the potential depending on the form of the function $f$ \cite{f(R),f(Phi R),f(Phii R)}. Likewise we get an MSTG when the original lagrangian is a more complicated function of multiple arguments of $R$, $\Box R$, $\nabla_\mu R \nabla^\mu R$, Gauss-Bonnet topological term, or Weyl tensor squared \cite{f(R G C etc}, as each such argument can contribute a scalar nonminimally coupled to $R$. (A function of arbitrary curvature invariants can be also turned into scalars, but the tensor part will not generally reduce to linear $R$ alone \cite{f(general R)}.) If the metric and connection are treated as independent variables, defining curvature scalars $R$ and $\mathcal{R}$, the resulting general hybrid metric-Palatini $f(R,\mathcal{R})$ gravity is equivalent to MSTG with two scalars \cite{general hybrid metric-Palatini}. A related construction called C-theory which continuously interpolates between metric and Palatini gravities, also possesses a biscalar-tensor representation \cite{C-theory}. In case the lagrangian is a function of higher derivatives of the curvature, $f(R, \Box^i R)$, each such argument can be converted to a nonmimimal scalar in MSTG \cite{f(box^i R)}. Moreover, a lagrangian of nonlocal gravity, characterized by derivatives in the inverse powers, $f(R, \Box^{-i} R)$, can be made local by again introducing auxiliary scalar fields nonminimally coupled to $R$ \cite{nonlocal,sasaki nonlocal}. The parametrized post-Newtonian (PPN) formalism is designed to describe slow motions in weak gravitational fields \cite{Will}, and can be utilized to confront the theory with high precision measurements in e.g.\ the Solar System. The original STG computation by Nordtvedt \cite{Nordtvedt1970}, which assumed that the potential (mass) of the scalar field vanishes, has been generalized to studies of higher order effects \cite{2PN} and for models with altered kinetic term or extra scalar-matter couplings \cite{PPN_nonstandard_STG_without_potential}. An important lesson learned in the STG case is that making the scalar field massive by the inclusion of the potential modifies the PPN parameters \cite{Olmo_Perivolaropoulos,meie2013,Scharer:2014kya,meie2014}, so that the theory becomes viable in a much larger domain (cf.\ also Refs.\ \cite{massive_BD}, \cite{STG_neutron_stars}). This has been especially relevant for understanding the PPN behavior of $f(R)$ gravity \cite{ppn_fR}, equivalent to a subclass of STG. Of course, a similar effect is also present in the generalized STG or Horndeski theory \cite{Horndeski_ppn}. Curiously enough, in teleparallel theories where the scalar field is nonminimally coupled to torsion, the PPN parameters coincide with the ones of general relativity \cite{teleparallel_ppn} unless a boundary term is introduced to the action \cite{teleparallel_boundary_ppn}. In the pioneering MSTG paper Damour and Esposito-Far\`ese derived the PPN parameters in the Einstein frame and assuming the potential vanishes \cite{DamourEF}. The effect of the potential was also not considered in the Jordan frame computation for the constant diagonal kinetic term \cite{Berkin:1993bt} and more recently for a generic kinetic term \cite{erik2014}. These results were generalized to arbitrary frame and scalar field parametrization by using the formalism of invariants \cite{meie2015}. The PPN parameters for C-theory have been determined by relying on the correspondence with a subclass of MSTG in the massless (vanishing potential) limit \cite{C-theory PPN}, while the PPN parameter $\gamma$ for specific nonlocal gravity models has been found using the biscalar representation \cite{Koivisto:2008dh} as well as independently of MSTG \cite{Conroy:2014eja}. The purpose of this paper is to calculate the PPN parameter $\gamma$ for a Jordan frame MSTG with generic kinetic terms and arbitrary potential (but without derivative couplings). In Sec.\ \ref{Sec_2} we recall the Jordan frame MSTG action in different parametrizations. In the process we define a covariant metric on the space of scalar fields and the vector of nonminimal coupling, these allow us to clarify the invariant notion of ghosts and the meaning of nonminimal coupling. Next in Sec.\ \ref{Sec_3} we carry out the PPN computation for a point mass source and find that the effective gravitational constant as well as the PPN parameter $\gamma$ in general depend on the distance from the source. Sec.\ \ref{Sec_4} uncovers the geometric picture underlying this result in terms of the eigenvectors of the mass matrix. Sec.\ \ref{Sec_5} draws rough experimental bounds on the two scalars case from the Cassini tracking experiment. Sec.\ \ref{Sec_6} illustrates how to apply our formalism for various interesting examples of MSTG: nonminimally coupled Higgs SU(2) doublet, general hybrid metric-Palatini gravity, linear ($\Box^{-1}$) nonlocal gravity, and quadratic ($\Box^{-2}$) nonlocal gravity. The last section \ref{Sec_7} provides a summary and outlook. Some more technical calculations are given in the appendices. Appendix \ref{App_A} discusses when the mass matrix can be diagonalized. Appendix \ref{App_B} deals with the boundary value problem and the determination of integration constants. Appendix \ref{App_C} addresses the cases when the mass matrix can not be brought into diagonal form and there are higher dimensional Jordan blocks.
\label{Sec_7} In this paper we considered a generic multiscalar-tensor gravity (MSTG) with arbitrary coupling functions and potential (but no derivative couplings) in the Jordan frame, and computed the post-Newtonian parameter $\gamma$ using a point mass as a source. In the single field case the result reproduces the earlier study \cite{meie2013}, where a massive nonminimal scalar is known to modify the effective gravitational constant $G_{\mathrm{eff}}$ and the PPN parameter $\gamma$ by a correction which falls off exponentially in distance. The same effect persists in the multiscalar case, while $G_{\mathrm{eff}}$ (\ref{G_eff}) and $\gamma$ (\ref{gengamma}) now depend in an intricate way not only on the masses but also on the alignment of the fields in the field space (\ref{Gamma_def}). To describe the geometry of the field space we found it useful to introduce a metric (\ref{F_ab}) and a vector (\ref{k def}) which describes the nonminimal coupling of the scalars to gravity (spacetime curvature). These objects transform covariantly under the reparametrization of the scalars, i.e.\ a change of the coordinates and the corresponding local basis in the field space. The PPN calculation lead to the mass matrix (\ref{mass_matrix}) whose eigenvectors (or generalized eigenvectors) form a special basis in the field space. A nice interpretation can be given if the field space metric is positive definite (there are no ghosts among the scalars), as in Eq.\ (\ref{Gamma_interpretation}) each massive field gives a contribution depending on the angle between the respective mass eigenvector and the overall vector of nonminimal coupling. The situation is perhaps easier to grasp in the case of just two fields in the Brans-Dicke like parametrization where the formula (\ref{gamma_BD_2fields}) for $\gamma$ can be used to plot constraints on the theory parameters from the Cassini tracking experiment, see Fig.\ \ref{gamma_bounds}. As expected, for massless scalars the bounds on the asymptotic value of the Brans-Dicke $\omega$ are high, while sufficiently large masses remove any bound on $\omega$. A very interesting scenario would arise when one massless scalar is accompanied by a rather massive scalar, since in that case the experimental constraints on $\omega$ will be also greatly reduced compared to a single massless nonminimal scalar (provided the alignment in the field space is sufficiently favorable). Our results are very general and can be utilized to test the viability of a multitude of models, either originally formulated as MSTG or shown to be equivalent to MSTG. As an illustration of the formalism we considered four relevant examples: nonminimally coupled Higgs SU(2) complex doublet, general hybrid metric-Palatini gravity, linear ($\Box^{-1}$) nonlocal gravity, and quadratic ($\Box^{-2}$) nonlocal gravity. In the cases where an earlier PPN result was available in the literature for these examples, it agreed with the application of our formulas for that specific model. The computations of our paper were carried out in the Jordan frame. However, in various contexts and for several applications it is useful to focus upon the Einstein frame instead. Therefore it remains a task for future to give the results also in the Einstein frame, or even better, rephrase them in terms of the formalism of invariants \cite{meie2015} which facilitates their easy implementation in any frame and parametrization. Our insights to the role of the geometry on the field space are most likely just a first glimpse, and the formalism of invariants seems a fruitful tool to clarify these issues as well. The estimation of the numerical bounds on the model parameters assuming a characteristic distance from a point mass can be a rather crude approximation for an actual astrophysical experiment. For a more realistic situation one should integrate over geodesics and invoke an extended source, as has been done for a single nonminimally coupled field case \cite{Deng:2016moh, Zhang:2016njn}. Finally, to fully test the MSTG models the PPN weak field arena must be complemented by research on the strong field regime as well \cite{Berti:2015itd}.
16
7
1607.02356
We compute the parametrized post-Newtonian parameter γ in the case of a static point source for multiscalar-tensor gravity with completely general nonderivative couplings and potential in the Jordan frame. Similarly to the single massive field case γ depends exponentially on the distance from the source and is determined by the length of a vector of nonminimal coupling in the space of scalar fields and its orientation relative to the mass eigenvectors. Using data from the Cassini tracking experiment, we estimate bounds on a general theory with two scalar fields. Our formalism can be utilized for a wide range of models, which we illustrate by applying it to nonminimally coupled Higgs SU(2) doublet, general hybrid metric-Palatini gravity, linear (□<SUP>-1</SUP> ) and quadratic (□-<SUP>2</SUP>) nonlocal gravity.
false
[ "scalar fields", "nonminimal coupling", "general hybrid metric-Palatini gravity", "completely general nonderivative couplings", "multiscalar-tensor gravity", "Jordan", "linear", "γ", "potential", "two scalar fields", "the single massive field case", "the mass eigenvectors", "the Jordan frame", "the parametrized post-Newtonian parameter γ", "a general theory", "a static point source", "quadratic (□-<SUP>2</SUP>) nonlocal gravity", "bounds", "Higgs SU(2) doublet", "models" ]
10.170788
1.076899
89
12473477
[ "Glampedakis, Kostas", "Lasky, Paul D." ]
2016MNRAS.463.2542G
[ "The freedom to choose neutron star magnetic field equilibria: Table 1." ]
13
[ "Departamento de Física, Universidad de Murcia, Murcia, E-30100, Spain; Theoretical Astrophysics, University of Tübingen, Auf der Morgenstelle 10, Tübingen, D-72076, Germany", "Monash Centre for Astrophysics, School of Physics and Astronomy, Monash University, VIC 3800, Australia" ]
[ "2017PhRvD..96j3012G", "2018MNRAS.473.4272K", "2019MNRAS.489.4261M", "2020MNRAS.495.1360S", "2021MNRAS.503..875L", "2021MNRAS.503.2764F", "2021MNRAS.506.2985K", "2022ASSL..465..245D", "2022ApJ...928...53W", "2022pas..conf..231S", "2023MNRAS.519.1080G", "2023MNRAS.523.4089S", "2023PhRvD.108h4006S" ]
[ "astronomy" ]
1
[ "stars: magnetic field", "stars: neutron", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1988ApJ...325..725M", "1992ApJ...395..240R", "2002A&A...393..949P", "2004PhRvD..69d3001P", "2007A&A...469..275B", "2007A&A...470..303P", "2007Ap&SS.308..181H", "2007LRR....10....1A", "2008JFM...607..235Y", "2008MNRAS.385..531H", "2008MNRAS.385.2080C", "2009A&A...499..557R", "2009MNRAS.397..763B", "2009MNRAS.397..913C", "2010PNAS..107.7147K", "2011A&A...532A..30K", "2011ApJ...735L..20L", "2011MNRAS.410..805G", "2011MNRAS.417.2288M", "2012ApJ...760....1C", "2012MNRAS.420.1263G", "2012MNRAS.424..482L", "2013MNRAS.428L..26G", "2013MNRAS.432.1245F", "2013MNRAS.433.2445A", "2013MNRAS.434..123V", "2013MNRAS.434.1658M", "2013MNRAS.435L..43C", "2013PhRvD..88j3005L", "2013PhRvL.110g1101L", "2014MNRAS.442L..90K", "2014PhRvL.112q1101G", "2015MNRAS.447.3278B", "2015MNRAS.450.1638G", "2015MNRAS.452.3246P", "2015SSRv..191..315M", "2016MNRAS.455.1489P", "2016PNAS..113.3944G" ]
[ "10.1093/mnras/stw2115", "10.48550/arXiv.1607.05576" ]
1607
1607.05576_arXiv.txt
\label{sec:intro} The structure of magnetic fields in the interior of neutron stars has undergone intense scrutiny in recent years. The endeavour has been mostly motivated by X-ray observations of strongly magnetised neutron stars. For example, bursts and giant flares occurring in magnetars are commonly believed to be powered by these objects' super-strong magnetic fields \citep[see][for reviews]{woods06,mereghetti15}. Other isolated neutron stars---colloquially known by the sobriquet `magnificent seven'---are kept warm by their evolving magnetic fields and emit intense thermal radiation \citep{haberl07}. In fact, the resulting theory-observations synergy appears to converge to a notion of a magnetic field-dominated evolutionary link between various sub-families of the neutron star population \citep{kaspi10,vigano13}. For the neutron star theorist wishing to delve into the physics of neutron star magnetic fields, a reasonable starting point is determining the nature of magnetic field equilibria. There is a significant corpus of recent work on that topic, both analytical ~\citep[e.g.,][]{haskell08, reisenegger09,ciolfi09,mastrano11,glampedakis12a,lasky13b,ciolfi13} and numerical~\citep[e.g.,][]{braithwaite07,braithwaite09, lasky11,kiuchi11,ciolfi12,lander12,lander12b,fujisawa13,palapanidis15,bucciantini15}. One important result that has emerged is that magnetohydrodynamic (MHD) equilibria are greatly influenced by the properties of matter or, in other words, by the equation of state (EOS). The key physics factor is that of \textit{barotropic} versus \textit{non-barotropic} matter, the latter allowing a much larger family of magnetic field equilibria \citep{reisenegger09}. A closely related factor is that of the imposed \textit{symmetries} in the system: for practical reasons most existing work assumes axisymmetric (2-D) magnetic equilibria; very little is known of non-axisymmetric (3-D) equilibria. \begin{table*} \begin{minipage}{135mm} \caption{Overview of results. For each type of equation of state, symmetry of the magnetic field and type of gravity (i.e., Newtonian/general relativistic), we calculate whether an arbitrary magnetic field can be balanced by the degrees of freedom in the fluid, and hence whether one is arbitrarily free to prescribe a magnetic field. The final column of the table lists the section of the paper in which the relevant calculations can be found.}\label{tab:overview} \begin{tabular}{rcccc} \hline Equation of state & Symmetry & Gravity & Free to prescribe? & Section\\ \hline\hline Single fluid, barotropic & axi- \& non-axisymmetric & Newtonian & \xmark & \ref{sec:baro}\\ Single fluid, non-barotropic & axisymmetric & Newtonian & \cmark & \ref{sec:axi}\\ Single fluid, non-barotropic & non-axisymmetric & Newtonian & \xmark & \ref{sec:nonaxi}\\ Single fluid, elastic crust & axi- \& non-axisymmetric & Newtonian & \cmark & \ref{sec:crust}\\ Multifluid, cold $npe$ matter & axi- \& non-axisymmetric & Newtonian & \xmark & \ref{sec:coldsf} \\ Multifluid, hot $npe\mu$ matter & axisymmetric & Newtonian & \cmark & \ref{sec:hotsfaxi}\\ Multifluid, hot $npe\mu$ matter & non-axisymmetric & Newtonian & \xmark & \ref{sec:hotsfnonaxi}\\ \hline Single fluid, barotropic & axisymmetric & general relativistic & \xmark & \ref{sec:baroGS}\\ Single fluid, non-barotropic & axisymmetric & general relativistic & \cmark & \ref{sec:noGS}\\ \hline \end{tabular} \end{minipage} \end{table*} This paper addresses a more `global' issue: \textit{to what extent MHD equilibria in neutron stars are arbitrary for given EOSs, with or without axisymmetry.} Here, `arbitrary' means that a magnetic field $\bB$ can be freely prescribed (assuming it obeys $\bnabla \cdot \bB=0$), knowing that the available fluid degrees of freedom can satisfy the equilibrium equations. In order to address this crucial issue we start from the simplest case of a single-fluid barotropic model and subsequently consider more realistic neutron star models with stratification, superfluidity, departure from chemical equilibrium, muons and entropy. For each case we study both axisymmetric and non-axisymmetric equilibria. In some cases we find that the magnetic field {\it cannot} be freely prescribed, and is required to solve a Grad-Shafranov differential equation (i.e. the MHD equilibrium equation for the magnetic scalar `stream functions' and the fluid parameters of the non-magnetic star) while in other cases the field can be completely `user-specified'. We also find that non-axisymmetry entails additional constraints for the magnetic field. A summary of our main results can be found below in Section~\ref{sec:overview}. The paper is set out as follows: in \S\ref{singlefluid} we develop the general formalism for a single fluid star, applying this specifically to barotropic and stratified matter in \S\ref{sec:baro} and \S\ref{sec:stratified}, respectively. We apply specific examples of axi- and non-axisymmetric magnetic field equilibria in \S\ref{sec:axi} and \S\ref{sec:nonaxi}, respectively. In \S\ref{sec:crust} we digress and discuss magnetic field equilibrium in neutron star crusts. In \S\ref{sec:2fluid} we generalise the formalism to multifluid neutron stars with superfluid/superconducting components. We first treat the case of cold $npe$ matter and subsequently move to the most realistic model considered in this paper, accounting for the presence of muons and a finite temperature. Finally, in \S\ref{sec:GR} we consider full general relativistic (GR) MHD equilibria with both barotropic and non-barotropic EOS. \subsection{Overview of results} \label{sec:overview} This section provides a compact summary of our main results, having in mind the `fast-track' reader who may not have a penchant for detailed calculations and numerous equations. Our findings are shown in Table~\ref{tab:overview}. We can categorise our results in the following way: \begin{enumerate} \item\label{list:one} assuming axisymmetry, we first recover the well-known results for single-fluid barotropic stars (the field is required to solve a Grad-Shafranov equation) and single-fluid non-barotropic stars (the field can be freely prescribed). Moving to multifluid neutron stars (which are always non-barotropic) with cold $npe$ matter, we find that the system effectively behaves as a barotrope with an accompanying Grad-Shafranov equation. The addition of muons and entropy in the previous model (hot $npe\mu$ matter) restores the complete freedom in prescribing an MHD equilibrium. \item\label{list:two} non-axisymmetric equilibria in the core are \textit{never} completely arbitrary, even when the magnetic field is not required to solve a Grad-Shafranov equation. This is due to the presence of an additional constraint equation between the azimuthal and non-azimuthal magnetic force components. \item the transition from Newtonian MHD to GRMHD does not produce any qualitative changes and the conclusions of \ref{list:one} and \ref{list:two} remain valid (this why we show just two cases of GRMHD equilibria). \item The magnetic equilibrium in the crust can be freely specified provided we allow for a strained crust and the associated elastic force. \end{enumerate}
\label{sec:conclusions} The allowed space of MHD equilibria in neutron stars is dependent on the nature of matter, namely the equation of state. In this paper, we have surveyed one's freedom to arbitrarily prescribe magnetic equilibria for different types of matter (i.e. number of distinct fluids, composition and entropy gradients, deviations from chemical equilibrium), degrees of symmetry (i.e. axisymmetry/non-axisymmetry) and types of gravity (i.e. Newtonian/general relativistic). This freedom depends on whether there are available fluid degrees of freedom to balance the magnetic force. Our results are summarised as follows (see Table \ref{tab:overview} for a bird's eye view summary): \begin{enumerate} \item Axisymmetric systems have a rich spectrum of arbitrariness with respect to MHD equilibria. We have found that the usual Grad-Shafranov equation is not only a property of simple barotropic stellar models. It can also control the MHD equilibrium in stratified matter provided the latter is multifluid (e.g. $npe$ matter with neutron superfluidity). However, the addition of muons and entropy (hot $npe\mu$ matter) nullifies the Grad-Shafranov equation, eventually leading to freely specifiable magnetic fields (as in the case of single-fluid non-barotropic systems). Among other things, this freedom implies an arbitrary relative strength between the poloidal and toroidal field components. \item In non-axisymmetric systems, the additional azimuthal components of the equations of motion prevent the magnetic field from being arbitrarily specified. The resulting constraint, at least for the case of non-superconducting matter, leads to a pair of Grad-Shafranov-like equations for the magnetic field's scalar degrees of freedom (i.e. the Euler potentials). \item The transition from Newtonian to general relativistic gravity does not alter the above conclusions (but increases the complexity of the various equilibrium equations). \end{enumerate} This paper has solely focused on MHD equilibria. The dynamical \textit{stability} of these equilibria is a completely different and much harder question, with obvious repercussions for their astrophysical relevance. Recent work \citep{lander12} suggests that barotropic equilibria are generically unstable, but this may not be true for more realistic non-barotropic systems since the buoyancy force emerging in stratified matter is known to enhance stability~\citep[e.g.,][]{akgun13}. Another avenue for instability could be provided by the interplay between rotation and magnetic field crust-core coupling during the initial spin down of newly formed neutron stars~\citep{glampedakis15}. Somewhat surprisingly, our work has some bearing on the nature of the Hall equilibrium in neutron star crusts. This equilibrium refers to the asymptotic $t \to +\infty$ state of the magnetic induction equation when the field is sourced by electron currents (this is the so-called electron-MHD) and is set to evolve due to the Hall term, i.e. $ \partial_t \bB = \bnabla \times (\mathbf{v}_\re \times \bB ) \propto \bnabla \times \{ (\mathbf{J}_\re \times \bB)/n_\re \}$. We can immediately deduce that, in axisymmetry, the condition for Hall equilibrium is identical to Eqn.~(\ref{curl1}), thus leading to the Grad-Shafranov equation (\ref{GradS}) with $\rho \to \rho_\re$~\citep{gourgouliatos14}. It is straightforward to generalise this result to the full non-axisymmetric case; we predict that the Hall equilibrium should be described by our 3-D Grad-Shafranov equation~(\ref{3DGSeq}). It will be interesting to test this prediction with the recently developed numerical framework for 3-D Hall evolution in neutron star crusts~\citep{gourgouliatos16}. We hope that this paper will serve as a stepping stone for modelling the next-generation MHD equilibria in realistic neutron stars.
16
7
1607.05576
Our ability to interpret and glean useful information from the large body of observations of strongly magnetized neutron stars rests largely on our theoretical understanding of magnetic field equilibria. We answer the following question: is one free to arbitrarily prescribe magnetic equilibria such that fluid degrees of freedom can balance the equilibrium equations? We examine this question for various models for neutron star matter; from the simplest single-fluid barotrope to more realistic non-barotropic multifluid models with superfluid/superconducting components, muons and entropy. We do this for both axi- and non-axisymmetric equilibria, and in Newtonian gravity and general relativity. We show that, in axisymmetry, the most realistic model allows complete freedom in choosing a magnetic field equilibrium whereas non-axisymmetric equilibria are never completely arbitrary.
false
[ "non-axisymmetric equilibria", "magnetic field equilibria", "magnetic equilibria", "general relativity", "a magnetic field equilibrium", "complete freedom", "more realistic non-barotropic multifluid models", "Newtonian gravity", "entropy", "freedom", "both axi- and non-axisymmetric equilibria", "various models", "the equilibrium equations", "neutron star matter", "muons", "Newtonian", "useful information", "such that fluid degrees", "observations", "strongly magnetized neutron stars" ]
5.351501
2.909595
69
12516965
[ "Ward, Jonathan T.", "Austermann, Jason", "Beall, James A.", "Choi, Steve K.", "Crowley, Kevin T.", "Devlin, Mark J.", "Duff, Shannon M.", "Gallardo, Patricio A.", "Henderson, Shawn W.", "Ho, Shuay-Pwu Patty", "Hilton, Gene", "Hubmayr, Johannes", "Khavari, Niloufar", "Klein, Jeffrey", "Koopman, Brian J.", "Li, Dale", "McMahon, Jeffrey", "Mumby, Grace", "Nati, Federico", "Niemack, Michael D.", "Page, Lyman A.", "Salatino, Maria", "Schillaci, Alessandro", "Schmitt, Benjamin L.", "Simon, Sara M.", "Staggs, Suzanne T.", "Thornton, Robert", "Ullom, Joel N.", "Vavagiakis, Eve M.", "Wollack, Edward J." ]
2016SPIE.9914E..37W
[ "Mechanical designs and development of TES bolometer detector arrays for the Advanced ACTPol experiment" ]
5
[ "Univ. of Pennsylvania (United States)", "NIST Quantum Devices Group (United States)", "NIST Quantum Devices Group (United States)", "Princeton Univ. (United States)", "Princeton Univ. (United States)", "Univ. of Pennsylvania (United States)", "NIST Quantum Devices Group (United States)", "Cornell Univ. (United States)", "Cornell Univ. (United States)", "Princeton Univ. (United States)", "National Institute of Standards and Technology (United States)", "NIST Quantum Devices Group (United States)", "Univ. of Pennsylvania (United States)", "Univ. of Pennsylvania (United States)", "Cornell Univ. (United States)", "SLAC National Accelerator Lab. (United States)", "Univ. of Michigan (United States)", "Univ. of Pennsylvania (United States)", "Univ. of Pennsylvania (United States)", "Cornell Univ. (United States)", "Princeton Univ. (United States)", "Princeton Univ. (United States)", "Pontificia Univ. Catolica de Chile (Chile)", "Univ. of Pennsylvania (United States)", "Princeton Univ. (United States)", "Princeton Univ. (United States)", "West Chester Univ. of Pennsylvania (United States)", "NIST Quantum Devices Group (United States)", "Cornell Univ. (United States)", "NASA Goddard Space Flight Ctr. (United States)" ]
[ "2018JLTP..193..288V", "2019BAAS...51c..37K", "2019BAAS...51c..51K", "2019arXiv190704473A", "2022SPIE12190E..1KH" ]
[ "astronomy", "physics" ]
5
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "2010ApJ...722.1148F" ]
[ "10.1117/12.2233746", "10.48550/arXiv.1607.05754" ]
1607
1607.05754_arXiv.txt
The Atacama Cosmology Telescope (ACT) \cite{ACT} is a six-meter diameter telescope located at an elevation of 5,190 meters on Cerro Toco in the Andes mountains of northern Chile. With a broader frequency range and higher sensitivity than the first polarimeter used for ACT (ACTPol\cite{BobInstrument}), AdvACT aims to make high resolution measurements of the polarized Cosmic Microwave Background (CMB) radiation over a range of angular scales. These observations will probe properties of the universe such as the tensor to scalar ratio $r$ and the sum of the neutrino masses, while also further constraining standard $\Lambda$CDM model parameters (e.g. Spergel et al. 2003\cite{Spergel} and Sievers, Hlozek, Nolta et al. 2013\cite{Sievers}). In order to achieve these and other science goals, the AdvACT instrument \cite{HendersonLTD} will consist of four multichroic, polarization-sensitive TES bolometer detector arrays\cite{LiLTD}. Each array will be sensitive to two frequency bands: 150 and 230 GHz in a high frequency (HF) array, 90 and 150 GHz in two middle frequency (MF) arrays, and 28 and 41 GHz in a low frequency (LF) array. The detectors are cryogenically cooled to 100~mK using a dilution refrigerator (DR), which allows for continuous cooling, increases detector sensitivity, and is also required for many of the superconducting components of the detectors and readout to operate. This paper discusses the design, fabrication and assembly of the AdvACT detector and feedhorn wafers as well as the mechanical array structure used to house all of the 100~mK electronic components. We also describe new structural components that have been built to allow AdvACT array packages to integrate with the existing ACTPol receiver.
16
7
1607.05754
The next generation Advanced ACTPol (AdvACT) experiment is currently underway and will consist of four Transition Edge Sensor (TES) bolometer arrays, with three operating together, totaling 5800 detectors on the sky. Building on experience gained with the ACTPol detector arrays, AdvACT will utilize various new technologies, including 150 mm detector wafers equipped with multichroic pixels, allowing for a more densely packed focal plane. Each set of detectors includes a feedhorn array of stacked silicon wafers which form a spline profile leading to each pixel. This is then followed by a waveguide interface plate, detector wafer, back short cavity plate, and backshort cap. Each array is housed in a custom designed structure manufactured from high purity copper and then gold plated. In addition to the detector array assembly, the array package also encloses cryogenic readout electronics. We present the full mechanical design of the AdvACT high frequency (HF) detector array package along with a detailed look at the detector array stack assemblies. This experiment will also make use of extensive hardware and software previously developed for ACT, which will be modified to incorporate the new AdvACT instruments. Therefore, we discuss the integration of all AdvACT arrays with pre-existing ACTPol infrastructure.
false
[ "detectors", "multichroic pixels", "cryogenic readout electronics", "the detector array stack assemblies", "HF) detector array package", "the detector array assembly", "AdvACT", "pre-existing ACTPol infrastructure", "various new technologies", "back short cavity plate", "the ACTPol detector arrays", "high purity copper", "stacked silicon wafers", "backshort cap", "all AdvACT arrays", "150 mm detector wafers", "the array package", "the new AdvACT instruments", "a feedhorn array", "the AdvACT high frequency" ]
10.666058
2.899588
68
12472651
[ "Sollima, A.", "Ferraro, F. R.", "Lovisi, L.", "Contenta, F.", "Vesperini, E.", "Origlia, L.", "Lapenna, E.", "Lanzoni, B.", "Mucciarelli, A.", "Dalessandro, E.", "Pallanca, C." ]
2016MNRAS.462.1937S
[ "Searching in the dark: the dark mass content of the Milky Way globular clusters NGC288 and NGC6218" ]
22
[ "INAF Osservatorio Astronomico di Bologna, via Ranzani 1, I-40127 Bologna, Italy; Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Department of Physics, University of Surrey, Guildford GU2 7XH, UK", "Department of Astronomy, Indiana University, Bloomington, IN 47405, USA", "INAF Osservatorio Astronomico di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy", "Dipartimento di Astronomia, Universitá di Bologna, via Ranzani 1, I-40127 Bologna, Italy" ]
[ "2016MmSAI..87..614S", "2017ApJ...849..108G", "2017JCAP...04..036D", "2017MNRAS.464.2930H", "2017MNRAS.467..412P", "2017MNRAS.467..524B", "2017PhRvD..96f3001G", "2018ApJ...860...50F", "2018MNRAS.473..909B", "2018RAA....18..126S", "2019MNRAS.483.1400H", "2019MNRAS.484L.114K", "2019MNRAS.485.5345A", "2019MNRAS.487..147C", "2020JCAP...08..010F", "2020MNRAS.491..272H", "2020MNRAS.492.3169B", "2021A&A...652A..54A", "2021MNRAS.508.2688G", "2022A&A...666A.121R", "2024MNRAS.527.7495W", "2024MNRAS.529..331D" ]
[ "astronomy" ]
2
[ "methods: data analysis", "techniques: radial velocities", "stars: kinematics and dynamics", "stars: luminosity function", "mass function", "stars: Population II", "globular clusters: individual: NGC6218", "NGC288", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
[ "1916MNRAS..76..572E", "1966AJ.....71...64K", "1969ApJ...158L.139S", "1975ApJ...199L.143C", "1978MNRAS.183..341W", "1979AJ.....84..752G", "1979PAZh....5...77O", "1984ApJ...277..470P", "1984sv...bookR....F", "1985AJ.....90.1027M", "1988Natur.336...48G", "1989ApJ...339..178R", "1991A&A...250..113M", "1992AJ....104.2104L", "1993Natur.364..421K", "1993Natur.364..423S", "1994AJ....107.1397P", "1994AJ....108.1414W", "1994PASP..106..250S", "1995A&A...303..742D", "1995AJ....109..209G", "1995AJ....109..218T", "1996AJ....112.1487H", "1996IAUS..174..303H", "1996MNRAS.280..498D", "1996MNRAS.283L..72N", "1997MNRAS.289..898V", "2001MNRAS.322..231K", "2002AJ....123.1509B", "2002MNRAS.332..901N", "2002MNRAS.333..400C", "2002Msngr.110....1P", "2003A&A...409..523R", "2003MNRAS.340..227B", "2003MNRAS.344.1000B", "2003gmbp.book.....H", "2004MNRAS.353..829M", "2005ApJ...619..258M", "2005ApJ...625..825L", "2005ApJS..161..304M", "2006A&A...449..161D", "2006AJ....131.1163S", "2006Natur.442..539M", "2007AJ....133.1658S", "2007AJ....134..376D", "2007MNRAS.380..781S", "2008AJ....135.2055A", "2008ApJ...673..226P", "2008ApJ...676..594K", "2008MNRAS.391..942B", "2009A&A...500..785K", "2009ARA&A..47..371T", "2009ApJ...690.1370M", "2009ApJ...694.1498M", "2009ApJ...698.2093D", "2009MNRAS.396.2051B", "2010AJ....139..476P", "2010ARA&A..48..339B", "2010ApJ...708..698D", "2010ApJ...710.1063V", "2010MNRAS.406.2732L", "2010MNRAS.407.2241K", "2011A&A...530A..31L", "2011AJ....142....8S", "2011ApJ...726..108T", "2011ApJ...741...72C", "2011ApJ...741L..12B", "2011MNRAS.410.2698G", "2012A&A...538A..18B", "2012A&A...540A..16M", "2012ApJ...755..156S", "2012MNRAS.424..545N", "2012MNRAS.427..167S", "2012Natur.490...71S", "2013A&A...552A..49L", "2013A&A...558A.117L", "2013ApJ...763L..15M", "2013ApJ...769..107L", "2013ApJ...774..151M", "2013ApJ...777...69C", "2013MNRAS.428.3648I", "2013MNRAS.430L..30S", "2013MNRAS.432.2779B", "2014A&A...566A..58K", "2014MNRAS.437..415D", "2014MNRAS.439.2459H", "2014Natur.506..171P", "2015A&A...584A..52L", "2015ApJ...800....9M", "2015ApJ...803...29W", "2015ApJ...804...53H", "2015ApJ...805...65T", "2015ApJ...812..149W", "2015MNRAS.446.1820N", "2015MNRAS.448L..94S", "2015MNRAS.449L.100C", "2015MNRAS.451.2185S", "2015MNRAS.453.3278W", "2015MNRAS.453.3918M", "2015MNRAS.454.1658K" ]
[ "10.1093/mnras/stw1779", "10.48550/arXiv.1607.05612" ]
1607
1607.05612_arXiv.txt
\label{intro_sec} The relative contribution of luminous and dark matter to the overall mass budget of stellar systems contains crucial information on their nature, origin and evolution. According to the $\Lambda CDM$ cosmological paradigm, the structures in the Universe formed at high redshift through a hierarchical assembly of small fragments of non-baryonic matter (White \& Rees 1978). Galaxies of all morphological types are expected to form within these fragments being nowdays embedded in dark matter (DM) halos. This evidence comes from the discrepancy between the mass estimated for these stellar systems using the kinematics of their stars and their luminosities as tracers. In particular, the mass-to-light (M/L) ratios measured for these stellar systems range from 5 (for dwarf elliptical galaxies) to $>$1000 (for ultra faint dwarf spheroidals; Tollerud et al. 2011) i.e. several times larger than those predicted by population synthesis models ($1.5<M/L<2.5$; Bruzual \& Charlot 2003). Globular clusters (GCs) appears to stand-out from this scenario. Indeed, at odds with other DM dominated stellar systems populating contiguous regions of the luminosity-effective radius plane (Tolstoy, Hill \& Tosi 2009), they have low M/L ratios consistent with the hypothesis they are deprived of DM (McLaughlin \& van der Marel 2005; Strader, Caldwell \& Seth 2011). This difference might suggest a different formation scenario for these stellar systems which could form in collapsing DM-free gas clouds (see e.g. Kruijssen 2015). For this reason the M/L ratios of GCs are often used as a reference for stellar population studies and to validate the prediction of population synthesis models. On the other hand, low-mass DM halos surrounding GCs progenitors are hypothesized by some model of GC formation (Peebles 1984). These halos could be later stripped by the tidal interaction with the host galaxy leaving only minimal imprints in the structural and kinematical properties of present day GCs (Mashchenko \& Sills 2005). Observational claims of the possible existence of DM dominated GCs in the giant elliptical NGC 5128 have been recently put forward by Taylor et al. (2015). Non-baryonic DM is not the only invisible matter contained in stellar systems. Indeed, the final outcome of the stellar evolution process of stars with different masses is represented by remnants (white dwarfs, neutron stars and black holes) whose luminosities are comparable or even several orders of magnitude smaller than those of the least luminous Main Sequence (MS) stars. The estimate of the mass enclosed in dark remnants in a GC is complicated by the interplay between stellar and dynamical evolution in these stellar systems and by the uncertainties in their formation process (for a comprehensive discussion see Heggie \& Hut 1996). Indeed, both neutron stars (NSs) and black holes (BHs) are the compact remnants of massive ($M>8~M_{\odot}$) stars after their explosion as SNe II. The off-center onset of the explosive mechanism can transmit to the remnant a velocity kick often exceeding the cluster escape speed thus leading to its ejection outside the cluster (Drukier 1996; Moody \& Sigurdsson 2009). Moreover, the large mass contrast of BHs with respect to the mean cluster mass lead to a quick collapse of the BH population forming a dynamically decoupled sub-system in the central part of the cluster (Spitzer 1969). Scattering between multiple BHs leads to a prompt ejection of these objects . For the above mentioned reasons it was initially suggested that GCs should be deprived of BHs (Kulkarni, Hut \& McMillan 1993; Sigurdsson \& Hernquist 1993). However recent observational studies (see e.g. Strader et al. 2012; Chomiuk et al. 2013; Miller-Jones et al. 2015) have provided evidence of the possible presence of stellar mass BHs in a number of GCs and several theoretical studies (see e.g. Breen \& Heggie 2013; Heggie \& Giersz 2014; Morscher et al. 2013, 2015; Sippel \& Hurley 2013) have found that GCs might indeed still host a non-negligible fraction of these compact objects. On the other hand, the large fraction of X-ray binaries and pulsars per unit mass in GCs (up to a factor 100 larger than that estimated in the Galactic field; Clark 1975; Grindlay \& Bailyn 1988) indicates that a certain number of these objects must be present in these stellar systems. White dwarfs (WDs) are the natural outcome of the evolution of low/intermediate mass ($M<8~M_{\odot}$) stars after the expulsion of their envelopes occurring at the end of their asympthotic giant branch phase. As a consequence of the long lifetimes of WD progenitors and the typical negative slope of the initial mass function (Kroupa 2001; Bastian, Covey \& Meyer 2010) the fraction of WDs steadly increases during the cluster lifetime making them a significant contributor to the mass budget of a GC in the last stages of its evolution (Vesperini \& Heggie 1997; Baumgardt \& Makino 2003). Because of the particular form of the initial-final mass relation of low-mass stars (Kalirai et al. 2008) the mass spectrum of these objects is expected to be peaked at $M\sim0.5~M_{\odot}$ i.e. only slightly larger than the present-day mean stellar mass ($\sim0.4~M_{\odot}$) and significantly smaller than the typical turn-off mass in a GC ($\sim0.8~M_{\odot}$). So, at odds with NSs and BHs occupying always the high tail of the mass distribution of stellar objects, stars evolving into WDs change their ranking in mass within the whole GC stellar population during their lifetimes. Consequently, two-body relaxation is expected to produce a progressive migration of WDs from the center, where their massive progenitors sunk, towards the outer regions where stars with smaller masses are preferentially located. Observational evidence of this phenomenon seems to be provided by the analysis of the radial distribution of WDs in different regions of the cooling sequence in 47 Tucanae (Heyl et al. 2015). The interaction of a GC with the tidal field of its host galaxy further complicate the prediction of the retention fraction of dark remnants. Indeed, the ever continuing injections of kinetic energy favors the evaporation of the kinematically hottest (mainly low-mass) stars, while massive objects (like NSs, BHs and massive WD progenitors) are preferentially retained. N-body simulations indicate that the actual fraction of retained remnants has deep implications in the dynamical evolution of GCs (Contenta, Varri \& Heggie 2015) and in their present-day M/L ratios (L{\"u}tzgendorf, Baumgardt \& Kruijssen 2013). In this context, many observational analyses aimed at investigating the dark content in GCs have been conducted in the past years by comparing the M/L ratios estimated from stars kinematics and those predicted by stellar population synthesis models (Richer \& Fahlman 1989; Meylan \& Mayor 1991; Leonard, Richer \& Fahlman 1992; Piatek et al. 1994; Dirsch \& Richtler 1995; Ibata et al. 2013; L{\"u}tzgendorf et al. 2013; Kamann et al. 2014). One of the main drawbacks of this approach resides in the choice of the uncertain parameters affecting the M/L like the present-day mass function and the retention fraction of dark remnants (Shanahan \& Gieles 2015). In Sollima et al. (2012) we determined the dynamical and luminous masses of a sample of six Galactic GCs by fitting simultaneously their luminosity functions and their line-of-sight (LOS) velocity dispersion profiles with multimass analytical models leaving the present-day mass function (MF) as a free parameter and making an assumption on the dark remnants retention fraction. From this study we found that the derived stellar masses were systematically smaller than the dynamical ones by $\sim 40\%$. Although many hypotheses were put forward, the most favored interpretation linked such a discrepancy to a fraction of retained dark remnants larger than expected. Unfortunately, the robustness of the obtained result relies on the ability of the adopted specific model in reproducing the actual degree of mass segregation of the analysed clusters. In particular, both the luminous and the dynamical masses of the best-fit model were constrained by observables measured in the cluster core, a region where mass segregation effects are maximized and where only a small fraction of the cluster mass is contained, and then extrapolated to the whole cluster (see Sollima et al. 2015). Moreover, the approach adopted in that work did not allow to obtain information on the radial distribution of the dark mass across the cluster, therefore complicating any interpretation on its nature. In this paper we present the result of a model-independent analysis of the most extensive photometric and spectroscopic datasets available in the literature for two Galactic GCs, namely NGC 288 and NGC 6218 (M12) with the aim of deriving their dark mass content and radial distribution. These objects are two well studied GCs located in the southern emisphere which are particularly suited for this kind of studies. Both clusters are indeed relatively close ($d<9$ kpc) and characterized by a small central projected density ($\Sigma_{V}>18~mag~arcsec^{-2}$; Harris 1996, 2010 edition) and for this reason it is possible to sample their luminosity function and velocity dispersion profiles through photometric and spectroscopic surveys even within their cores with no significant crowding problems. In Sect. 2 we describe the observational dataset used in the analysis and the data reduction techniques. In Sect. 3 the methods to derive luminous and dynamical mass profiles for the two target clusters are described. The fractions of dark mass as a function of the projected distance from the clusters' centers are shown in Sect. 4. The comparison with the prediction of N-body simulations and analytical models is performed in Sect. 5 to quantify the expected effect of a sizable population of dark remnants. We summarize and discuss our results in Sect. 6.
The main result obtained in this paper is that the largest fraction ($>60\%$) of the mass content of both the analysed GCs is dominated by centrally concentrated non-luminous mass. The results presented here confirm what already found in Sollima et al. (2012). However, a significant improvement with respect to that work has been made here since: {\it i)} the adopted dataset of radial velocities is $\sim$5 times larger within the half-light radii of both clusters, {\it ii)} the analysis presented here is model-independent as it requires no assumptions on the degree of mass segregation among mass groups, and {\it iii)} in the present work we are able to determine the radial distribution of the dark mass. The obtained result seems not to be due to the heating produced by tidal effects or binary stars being therefore linked to a real excess of dark mass in the central region of these clusters. The most likely hypothesis is that such a large mass excess is produced by dark remnants sunk in the cluster core during the cluster lifetime as a result of mass segregation. In fact, the comparison with N-body simulations indicates that GCs losing a significant ($>70\%$) fraction of their initial mass contain a centrally concentrated core of dark remnants constituted mainly of WDs accounting for $\sim$50\% of the total cluster mass within the same radial extent explored in the present analysis. This value is slightly lower, but within the uncertainties, than that estimated in our analysis. It is worth noting that the N-body simulations considered in this work were not specifically run to reproduce the present-day structure of the two analysed clusters and are subject to a different tidal stress. So, only qualitative conclusions can be drawn from such a comparison. However, similar results are obtained by other authors: Giersz \& Heggie (2011) used Monte Carlo simulations with $2\times10^{6}$ particles to model the present-day structure of 47 Tuc and found that the fraction of WDs is $\sim51\%$ in the core and even larger in a scaled simulation run with a flatter initial mass function. Sippel et al. (2012) and L{\"u}tzgendorf et al. (2013) analysed a set of N-body simulations including natal kicks for NSs and BHs and found that after 12 Gyr $\sim40\%$ of the cluster mass is contained in WDs (and a larger fraction is expected in the core). According to Vesperini \& Heggie (1997) and Baumgardt \& Makino (2003) the fraction of WDs and NSs steadly increases as a cluster evolves and loses mass, becoming dominant when the cluster lost $>60\%$ of its stars. In this regard, a large mass-loss rate in the two analysed clusters is also suggested by their relatively flat mass functions (de Marchi, Pulone \& Paresce 2006; Paust et al. 2010; Sollima et al. 2012). On the basis of the above consideration, it is still possible that the observed overabundance of non-luminous mass in the two GCs analysed here is entirely due to a compact population of centrally segregated WDs. A valuable test to this hypothesis would be to measure the number of WDs in the central region of these clusters through photometric analyses. Unfortunately, to perform a complete census of WDs in these clusters it would be necessary to reach extremely faint magnitudes ($V\sim30-31$) with a good level of completeness, an unfeasible task even with HST. The possibile presence of an IMBH in the center of these two GCs, as suggested for many other Galactic GCs by L{\"u}tzgendorf et al. (2013; but see van der Marel \& Anderson 2010 and Lanzoni et al. 2013), cannot be excluded. In particular, according to the actual fraction of retained remnants, the mass of such an hypothetical BH could lie in the range $M_{IMBH}<7\times10^{3} M_{\odot}$. On the other hand, we believe that given the above apparent degeneracy and the large uncertainty in the expected dark remnants fractions, any claim of IMBH detection is inappropriate. Similarly, the presence of an IMBH can significantly affect the determination of the actual fraction of remnants. Another possibility is that the mass excess observed here is due to a modest amount of non baryonic DM. Previous investigations in this regard made on the outer halo GC NGC 2419 seem to rule out the presence of a significant amount of DM within the stellar extent of this cluster (Baumgardt et al. 2009; Conroy, Loeb \& Spergel 2011) although the involved uncertainties in the anisotropy profile leave some room for a small DM content (Ibata et al. 2013). Note that in this scenario the DM halo should extend far beyond the extent of the stellar component. N-body simulations assuming GCs surrounded by cored DM halos predict that, because of the interaction with the cluster stars (Baumgardt \& Mieske 2008) and stripping by the tidal field of the host galaxy, the central parts of GCs might be left relatively poor in DM because at the present time the DM either populates the outer region of the cluster or has been stripped (Mashchenko \& Sills 2005; Pe{\~n}arrubia, Navarro \& McConnachie 2008). In this scenario, the high concentration of dark mass evidenced in our analysis would favor a cuspy shape of the surviving DM halo dominating the cluster mass budget in the central region of these clusters. This is however in contrast with theoretical considerations (where interactions with baryons are expected to remove DM cusps; Navarro, Eke \& Frenk 1996; Mo \& Mao 2004; Mashchenko, Couchman \& Wadsley 2006; Del Popolo 2009; Di Cintio et al. 2014; Pontzen \& Governato 2014; Nipoti \& Binney 2015). Because of the limited sample of GCs analysed here it is not possible to check the presence of correlation between dark mass content and other dynamical and general parameters. Future studies addressed to the extension of this analysis to a larger sample of GCs will help to clarify the nature of the dark mass and the effect of the various dynamical processes occurring during the cluster evolution on its fraction.
16
7
1607.05612
We present an observational estimate of the fraction and distribution of dark mass in the innermost region of the two Galactic globular clusters NGC 6218 (M12) and NGC 288. Such an assessment has been made by comparing the dynamical and luminous mass profiles derived from an accurate analysis of the most extensive spectroscopic and photometric surveys performed on these stellar systems. We find that non-luminous matter constitutes more than 60 per cent of the total mass in the region probed by our data (R &lt; 1.6 arcmin ∼ r<SUB>h</SUB>) in both clusters. We have carefully analysed the effects of binaries and tidal heating on our estimate and ruled out the possibility that our result is a spurious consequence of these effects. The dark component appears to be more concentrated than the most massive stars suggesting that it is likely composed of dark remnants segregated in the cluster core.
false
[ "M12", "NGC 288", "dark mass", "tidal heating", "dark remnants", "non-luminous matter", "the cluster core", "lt", "binaries", "R", "Galactic", "these stellar systems", "both clusters", "the two Galactic globular clusters", "distribution", "a spurious consequence", "the total mass", "the innermost region", "the effects", "these effects" ]
9.922097
7.502296
-1
12593492
[ "Desmond, Harry" ]
2017MNRAS.464.4160D
[ "A statistical investigation of the mass discrepancy-acceleration relation" ]
86
[ "Kavli Institute for Particle Astrophysics and Cosmology and Physics Department, Stanford University, Stanford, CA 94305, USA; SLAC National Accelerator Laboratory, Menlo Park, CA 94025, USA" ]
[ "2016PhRvL.117t1101M", "2016arXiv160906642M", "2017A&A...603A..11P", "2017A&A...603A..65T", "2017A&A...607A.108H", "2017ApJ...836..152L", "2017ApJ...837..179D", "2017IJMPD..2650118C", "2017MNRAS.469.1630H", "2017MNRAS.471.1841N", "2017MNRAS.471L..11D", "2017MNRAS.472..765T", "2017MNRAS.472L..35D", "2017PhRvL.118p1103L", "2017arXiv171003096P", "2018A&A...615A...3L", "2018ApJ...863..107D", "2018JCAP...03..038F", "2018MNRAS.473.4033B", "2018MNRAS.474.3125T", "2018MNRAS.474.3152D", "2018MNRAS.477.4768B", "2018MNRAS.480.2660B", "2018arXiv180404484M", "2018arXiv180509207B", "2018arXiv180810545B", "2019A&A...623A.123G", "2019ApJ...877...18C", "2019ApJ...882...46W", "2019ApJ...885...87M", "2019MNRAS.483L..98K", "2019MNRAS.484..239D", "2019MNRAS.485.1886D", "2019MNRAS.487.1653B", "2019MNRAS.487.2148G", "2019MNRAS.487.5291B", "2019MNRAS.488L..41T", "2019MNRAS.489..771N", "2019MNRAS.489.1805H", "2019PhRvX...9c1020R", "2020ApJ...890..173W", "2020ApJ...896...70T", "2020ApJ...904...51C", "2020ApJ...905..135B", "2020IAUS..353..144M", "2020MNRAS.492.1671Z", "2020MNRAS.492.2698M", "2020MNRAS.492.5865C", "2020MNRAS.494.4015W", "2020MNRAS.495.3974B", "2020MNRAS.496.1077P", "2020MNRAS.497L..62B", "2020MNRAS.498.5885B", "2020PDU....2800478C", "2020PhRvD.101h4015I", "2021A&A...650A.113B", "2021ApJ...923...68O", "2021MNRAS.505.4555Z", "2021MNRAS.506.3205S", "2021MNRAS.507..632P", "2021PDU....3100765P", "2021PDU....3300854P", "2022ApJ...927..198L", "2022ApJ...929...48G", "2022ApJ...941...55C", "2022CQGra..39g5001P", "2022MNRAS.509.2800B", "2022MNRAS.514.4026S", "2022MNRAS.517..130P", "2022NatAs...6...35L", "2022PASJ...74.1441C", "2022Symm...14.1331B", "2022arXiv220710638A", "2023A&A...674A.209S", "2023ApJ...952..105G", "2023JCAP...04..048D", "2023MNRAS.518..257M", "2023MNRAS.521.1817D", "2023MNRAS.522.4003P", "2023MNRAS.525.6130S", "2023MNRAS.526.3342D", "2023MNRAS.526.5861Y", "2023arXiv230519289L", "2024MNRAS.530.1349M", "2024MNRAS.530.1781D", "2024arXiv240510019M" ]
[ "astronomy" ]
5
[ "galaxies: formation", "galaxies: fundamental parameters", "galaxies: haloes", "galaxies: kinematics and dynamics", "dark matter", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stw2571", "10.48550/arXiv.1607.01800" ]
1607
1607.01800_arXiv.txt
\label{sec:intro} The internal motions of galaxies are largely set by dark matter, which outweighs baryonic matter by at least five to one. A key goal of galaxy astrophysics is to relate the visible and dark mass in any given system, which to first order means determining the correlations between the structural parameters of galaxies (e.g. $M_*$, $M_\mathrm{gas}$, $R_\mathrm{d}$ and Hubble type $T$) and those of dark matter haloes (e.g. $M_\mathrm{vir}$, $c$, and $\lambda$). This programme requires extensive observation of the rotation and velocity dispersion profiles of galaxies and their baryonic mass distributions, combined with detailed dynamical modelling under a variety of assumptions about the galaxy--halo connection. Traditional summaries of the relation between the baryonic mass distribution and internal motion of galaxies relate one-point statistics of these functions, reducing the former to a total galaxy mass and size, and the latter to a single measure of velocity. These are the Tully--Fisher, mass--size and Faber--Jackson relations, and Fundamental Plane. More information, however, can be found in the full \emph{radial} interdependence of mass and velocity, and analysis of this may be expected to afford not only a more stringent test of specific galaxy formation models, but also a richer foundation for a bottom-up determination of the galaxy--halo connection. The local correlation of dark and visible matter is usefully described by the ratio of enclosed dynamical mass (determined kinematically) to enclosed baryonic mass (determined photometrically). A proxy for this quantity is $V^2_\mathrm{tot}(r)/V^2_\mathrm{bar}(r)$, to which it is equal in the case of spherical mass distribution. $V^2_\mathrm{tot}(r)/V^2_\mathrm{bar}(r)$ is known as the ``mass discrepancy''~\citep{MG99}, and will be denoted by $\mathcal{D}$ hereafter. In general, $\mathcal{D}$ may be any function of position $r$ within a galaxy, with parametric dependence on global galaxy properties $X$ (e.g. $M_*$, $M_\mathrm{gas}$, $R_\mathrm{d}$, $T$). The full functional form of the mass discrepancy, $\mathcal{D}(r; X)$, is a kinematic parametrization of the relation between dark matter and baryons, and a fortiori of the galaxy--halo connection. While in principle $\mathcal{D}$ may have an arbitrary dependence on position, its utility is significantly enhanced by the fact that it is known to correlate strongly with acceleration $a(r)$ in all galaxies in which it has been measured in detail~\citep{Sanders_MDAcc, MG99, McGaugh_MDAcc, Tiret_Combes, Famaey_McGaugh, Janz}. This allows us to focus our attention on $\mathcal{D}(a(r); X)$, which is known as the ``mass discrepancy--acceleration relation'' or MDAR. The information content of the MDAR relative to one-point summaries may be approximately established by a simple counting argument: while the latter are limited to one data point per galaxy, the former contains as many data points as one can measure across a galaxy's entire rotation curve. The aim of this paper is to construct a framework for using the MDAR to test galaxy formation models, and hence deduce the dynamically relevant correlations of the galaxy--halo connection. Although the MDAR has been known observationally for decades, few studies have sought to systematically extract the information contained within it. The tightness of the relation is used by some to argue against all $\Lambda$CDM-based models of galaxy formation (e.g.~\citealt{MG98, MG99, McGaugh_MDAcc, Kroupa_Falsification, TTP, Kroupa}), where stochasticity in the galaxy--halo connection may be expected to introduce significant scatter into the relation between $\mathcal{D}$ and $a$. In addition, it is argued that the presence of a ``characteristic acceleration'' $a \approx 10^{-10} \: \mathrm{m\:s}^{-2}$ at which $\mathcal{D}$ consistently becomes $\sim 1$ (indicating a dynamically insignificant quantity of dark matter) is not compatible with standard theory. A generic galaxy--halo connection would predict a spread in $\mathcal{D}$ at high $a$ and no clear transition in acceleration space between the dark matter and baryon-dominated regimes. Other authors, however, claim that the salient features of the MDAR arise naturally in $\Lambda$CDM models that have been tuned to match the Tully--Fisher~\citep{vdB} or $M_*-M_\mathrm{halo}$, $M_*-R_\mathrm{d}$, and $M_*-M_\mathrm{gas}$ relations~\citep{DC}, and consensus concerning the relation's significance does not seem at hand. No study to date has systematically investigated the dependence of $\mathcal{D}$ on global galaxy properties ($X$) at fixed $a$, or quantified the correlation of MDAR residuals with $r$. Our specific task is twofold. First, we create a set of statistics to quantify four significant features of the MDAR: its shape, its scatter, the presence of a ``characteristic acceleration,'' and the correlation of its residuals with other variables ($M_*$, $M_\mathrm{gas}$, $R_\mathrm{d}$, $T$ and $r$). This is motivated in part by the prevalence of largely qualitative claims in the literature concerning the compatibility of the MDAR with various models, from which it may be difficult to determine the exact degree or nature of the agreement. Our statistics enable the conversion of verbal assertions into precise statistical comparisons, which we hope will sharpen discussion of the MDAR regardless of theoretical perspective. The salience of these statistical features, however, is best appreciated in the context of specific model expectations. Our second task, therefore, is to develop a semi-empirical framework in $\Lambda$CDM to generate predictions for the MDAR. We adopt a fully bottom-up methodology, beginning with the simplest and best motivated correlations between galaxy and halo variables, and introducing more when required by the data. By comparing the predicted and observed MDARs, we will deduce the extent to which semi-empirical models are able to account for the significant aspects of the relation, and the concrete extensions to basic models that are required to match the relation's more detailed features. We intend in this way to lay the groundwork for a phenomenological determination of the galaxy--halo connection from information-rich relations such as the MDAR, as well as formulate precise tests for specific models. The starting point of our framework is the technique of halo abundance matching (AM), which imposes a nearly monotonic relationship between galaxy stellar mass and halo mass or velocity at a particular epoch~\citep{Kravtsov,Conroy,Behroozi_2010,Guo,Moster}. From a phenomenological perspective, AM specifies the relation between stellar mass and halo mass and concentration required to fit clustering~\citep{Reddick, Lehmann} and dynamical~\citep{DW15, DW16} observations, but in its basic form neglects gas mass as well as galaxy size and type. We therefore augment the model by allowing correlations of these variables with $M_\mathrm{vir}$ and $c$ at fixed $M_*$. For a given set of correlations constituting the galaxy--halo connection, we generate a large number of mock data sets from our theoretical population with baryonic properties identical to the real data and halo properties specified by the model. We then calculate the MDAR statistics of these mock data sets, and compare with the observations. The MDAR is considered by some an important piece of evidence in favour of Modified Newtonian Dynamics (MOND) as an alternative to $\Lambda$CDM for solving the missing mass problem, and is a central relation in MOND phenomenology (see~\citealt{Famaey_McGaugh} and references therein). Indeed, the founding papers of MOND were the first to predict that $\mathcal{D}$ would be more tightly correlated with acceleration than velocity or galactocentric radius~\citep{Milgrom1, Milgrom2, Milgrom3}, a hypothesis not verified empirically for many years. In MOND, the MDAR is a direct manifestation of the breakdown of Newtonian gravity or mechanics at low acceleration $a < a_0 \approx 1.2 \times 10^{-10} \: \mathrm{m\:s}^{-2}$, with the result that it is predicted to have negligible intrinsic scatter, residuals systematically uncorrelated with any other variable, and a clear acceleration scale $a=a_0$. While the MOND MDAR must go to 1 at $a \gg a_0$ (the Newtonian regime) and has a shape fixed by the theory at $a \ll a_0$ (the deep-MOND regime), the behaviour at intermediate $a$ is specified by an ad hoc interpolating function. Our statistical analysis will shed light on the compatibility of the observations with the MOND hypothesis, and our comparison with $\Lambda$CDM models will bring into sharper focus the relative evidence accorded to MOND and $\Lambda$CDM by the relation. The structure of this paper is as follows. In Section~\ref{sec:data} we describe our observational MDAR sample and the $N$-body simulations on which we build our theoretical framework. In Section~\ref{sec:method} we lay out our procedure for constructing the galaxy--halo connection and document the MDAR statistics with which we evaluate our models. In Section~\ref{sec:results} we present our comparison of theory and data, first for a fiducial model best motivated by prior analyses, and then allowing variations in our model assumptions to maximise agreement with the MDAR. Section~\ref{sec:discussion} discusses literature studies in the context of our results, locates our parameter constraints and model requirements relative to previous findings, and elaborates on the broader implications for our understanding of galaxy astrophysics. Section~\ref{sec:conclusion} concludes.
\label{sec:conclusion} The MDAR provides a map between the distribution of a galaxy's baryonic and dark matter, and therefore contains crucial information about the galaxy--halo connection. In this paper we have laid the groundwork for extracting this information for use in evaluating models of galaxy formation. We have analysed the MDAR using a set of 16 statistics that quantify its four most important features: its shape, its scatter, the presence of a ``characteristic acceleration scale'' beyond which mass discrepancy consistently goes to $\sim 1$, and the correlation of its residuals with other galaxy properties. In addition to using these statistics to focus discussion of the observed relation itself, we have engaged them to construct a data-driven framework for the galaxy--halo connection. Building upwards from the simplest case of stellar mass-based abundance matching in $\Lambda$CDM, we have successively incorporated selection effects, a correlation between galaxy size and halo concentration, and a mass-dependent prescription for the impact of disc formation on halo density profiles. Comparing to data from the \textsc{sparc} sample, our most significant findings are as follows. \begin{itemize} \item{} A basic AM model readily accounts for several features of the MDAR, including its approximate overall shape, its normalisation and scatter at high acceleration, and the independence of its residuals on stellar and gas mass. \item{} Nevertheless, the predicted MDAR has significantly too high a normalisation and scatter at low acceleration, and too high a scatter in an averaged sense over the whole relation. This remains true under highly conservative assumptions for galaxy formation, including no scatter in the galaxy--halo connection and a significant reduction in the spread of halo concentrations associated with galaxies of given $M_*$ and $R_\mathrm{d}$. This indicates too much dark matter mass predicted in the outer regions of high surface brightness galaxies and (especially) in low surface brightness galaxies, and too large a spread therein. In addition, dark matter is more concentrated towards the centres of galaxies than the MDAR suggests. These discrepancies argue for halo expansion in response to disc formation, and a quantitative resolution requires also the exclusion of a large fraction ($\sim 50$ per cent) of the haloes with highest concentration at each stellar mass. \item{} We devise six statistics to capture aspects of the ``characteristic acceleration scale'' the MDAR is sometimes said to exhibit, describing its acceleration behaviour -- and that of individual galaxies within it -- at low mass discrepancy. Although our models cannot simultaneously reproduce the observed values of each of these statistics, we find no grounds for the claim that the transition region between baryonic and dark matter domination is sharper than expected by standard galaxy formation in $\Lambda$CDM. \item{} The MDAR may be used to detect correlations of halo properties with (at least) three galaxy properties at fixed stellar mass: disc size, Hubble type, and gas mass. Our analysis provides weak evidence for an anticorrelation of halo mass or concentration with galaxy size and type at fixed stellar mass ($\sim 2.3 \sigma$ and $\sim 1.7 \sigma$ respectively), but no evidence for such a correlation with gas mass. \end{itemize} We hope that this work will stimulate interest in the MDAR as a source of information about the galaxy--halo connection. Looking forward, we identify three ways in which further progress could be made. First, additional models and assumptions need to be tested against the relation to the level of rigour achieved here. It is unclear to what extent the outputs of many hydrodynamical simulations or semi-analytic models, for example, are consistent with MDAR statistics. Specific implementations of the correlations of halo properties with galaxy variables besides $M_*$ (e.g. AM using total baryonic mass, or the angular momentum partition model of~\citealt{MMW}) should be tested individually. Second, a firmer theoretical basis needs to be given for the empirical galaxy--halo correlations argued for by the MDAR; this may be possible by mapping hydrodynamical simulations onto phenomenological frameworks such as ours. Finally, an increase in the size and precision of MDAR data sets may be expected to yield a considerable gain in the constraining power of analyses of this type, pinning down the values of the statistics in the real data and reducing the widths of their distributions in the mock data. The full power of the MDAR likely remains to be harnessed.
16
7
1607.01800
We use the mass discrepancy-acceleration relation (the correlation between the ratio of total-to-visible mass and acceleration in galaxies; MDAR) to test the galaxy-halo connection. We analyse the MDAR using a set of 16 statistics that quantify its four most important features: shape, scatter, the presence of a `characteristic acceleration scale', and the correlation of its residuals with other galaxy properties. We construct an empirical framework for the galaxy-halo connection in LCDM to generate predictions for these statistics, starting with conventional correlations (halo abundance matching; AM) and introducing more where required. Comparing to the SPARC data, we find that: (1) the approximate shape of the MDAR is readily reproduced by AM, and there is no evidence that the acceleration at which dark matter becomes negligible has less spread in the data than in AM mocks; (2) even under conservative assumptions, AM significantly overpredicts the scatter in the relation and its normalization at low acceleration, and furthermore positions dark matter too close to galaxies' centres on average; (3) the MDAR affords 2σ evidence for an anticorrelation of galaxy size and Hubble type with halo mass or concentration at fixed stellar mass. Our analysis lays the groundwork for a bottom-up determination of the galaxy-halo connection from relations such as the MDAR, provides concrete statistical tests for specific galaxy formation models, and brings into sharper focus the relative evidence accorded by galaxy kinematics to LCDM and modified gravity alternatives.
false
[ "halo mass", "galaxy kinematics", "galaxy size", "other galaxy properties", "specific galaxy formation models", "galaxies", "fixed stellar mass", "low acceleration", "acceleration", "halo abundance matching", "AM mocks", "MDAR", "AM", "conventional correlations", "relations", "concrete statistical tests", "sharper focus", "Hubble type", "LCDM", "dark matter" ]
10.165664
2.18853
51
5251088
[ "Fish, Vincent L.", "Akiyama, Kazunori", "Bouman, Katherine L.", "Chael, Andrew A.", "Johnson, Michael D.", "Doeleman, Sheperd S.", "Blackburn, Lindy", "Wardle, John F. C.", "Freeman, William T.", "the Event Horizon Telescope Collaboration" ]
2016Galax...4...54F
[ "Observing—and Imaging—Active Galactic Nuclei with the Event Horizon Telescope" ]
60
[ "Massachusetts Institute of Technology, Haystack Observatory, Westford, MA 01886, USA", "Massachusetts Institute of Technology, Haystack Observatory, Westford, MA 01886, USA; Japan Society for the Promotion of Science, Chiyoda (Tokyo) 102-0083, Japan", "Massachusetts Institute of Technology, Computer Science and Artificial Intelligence Laboratory, Cambridge, MA 02139, USA", "Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA", "Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA", "Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA", "Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA", "Physics Department, Brandeis University, Waltham, MA 02453, USA", "Massachusetts Institute of Technology, Computer Science and Artificial Intelligence Laboratory, Cambridge, MA 02139, USA; Google Inc., 1600 Amphitheatre Parkway, Mountain View, CA 94043, USA", "-" ]
[ "2017A&A...601A..52B", "2017A&ARv..25....2P", "2017A&ARv..25....4B", "2017ApJ...850..172J", "2017CaJPh..95.1299M", "2017JETPL.106..637D", "2017JPhCS.934a2044D", "2017arXiv170504776A", "2018A&A...613A...2B", "2018AJ....156...15M", "2018PhLB..781..651D", "2018PhRvD..97f4021T", "2018PhRvD..97f4041H", "2018PhRvD..97l4024J", "2018arXiv181010657B", "2018arXiv181100047W", "2019ApJ...870....6Z", "2019ApJS..241...33L", "2019GReGr..51...81D", "2019GReGr..51..137P", "2019IJMPD..2841005D", "2019JCAP...03..046W", "2019JETP..128..578D", "2019PhRvD..99b4035P", "2019PhRvD..99d1303W", "2019PhRvD..99d3002G", "2019PhRvD.100b4020V", "2019PhRvD.100b4055G", "2019PhRvD.100d4057B", "2019arXiv191000013A", "2020AdSpR..65..712B", "2020CQGra..37h7001V", "2020MNRAS.499.1561Z", "2020PhRvD.101d1301B", "2020PhRvD.102d4038K", "2020PhyU...63..583D", "2020Univ....6..154D", "2021ApJ...923..246G", "2021ChPhC..45a5105L", "2021JCAP...09..037B", "2021PhRvD.103b4017K", "2021PhRvD.103d4046B", "2022JCAP...02..034L", "2022JCAP...05..020B", "2022JCAP...09..066B", "2022PhRvD.105f4073B", "2022PhRvD.105h3002R", "2022PhRvD.106b4039S", "2022PhRvD.106j4024D", "2022arXiv220402026S", "2022arXiv220706034B", "2023CQGra..40p5007V", "2023EPJC...83..171B", "2023JCAP...02..030A", "2023JCAP...11..099R", "2023PhRvD.107d4031A", "2023PhRvD.107l4003S", "2024IAUS..380..275M", "2024MNRAS.528..735A", "2024arXiv240403755D" ]
[ "astronomy" ]
12
[ "galaxies: jets", "galaxy: center", "techniques: high angular resolution", "techniques: image processing", "techniques: interferometric", "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.3390/galaxies4040054", "10.48550/arXiv.1607.03034" ]
1607
1607.03034_arXiv.txt
The Event Horizon Telescope has been able to probe Sgr~A*, M87, and other AGN sources at very high angular resolution, addressing fundamental physical and astrophysical questions associated with accreting black holes. Improvements in sensitivity and baseline coverage, notably the inclusion of ALMA in 2017, will significantly increase the capacity of the EHT to produce images. New imaging techniques are being developed to use robust VLBI observables to make the most of the relatively sparse baseline coverage of the EHT. These techniques have broad applicability beyond the EHT (e.g., on longer-wavelength VLBI observations). Validation of these algorithms on real longer-wavelength VLBI data demonstrate their potential for reconstructed images with greater resolution and fidelity than currently provided by CLEAN. Observers using data from longer-wavelength VLBI arrays, such as the VLBA and GMVA, may find these methods useful. \vspace{6pt}
16
7
1607.03034
Originally developed to image the shadow region of the central black hole in Sagittarius A* and in the nearby galaxy M87, the Event Horizon Telescope (EHT) provides deep, very high angular resolution data on other active galactic nucleus (AGN) sources too. The challenges of working with EHT data have spurred the development of new image reconstruction algorithms. This work briefly reviews the status of the EHT and its utility for observing AGN sources, with emphasis on novel imaging techniques that offer the promise of better reconstructions at 1.3 mm and other wavelengths.
false
[ "other active galactic nucleus", "other wavelengths", "AGN sources", "novel imaging techniques", "better reconstructions", "EHT data", "sources", "AGN", "emphasis", "EHT", "Sagittarius A", "M87", "the Event Horizon Telescope", "deep, very high angular resolution data", "the central black hole", "the nearby galaxy", "the promise", "1.3 mm", "the EHT", "the shadow region" ]
8.918159
3.680691
-1
12619768
[ "Brdar, Vedran", "Kopp, Joachim", "Liu, Jia" ]
2017PhRvD..95e5031B
[ "Dark gamma-ray bursts" ]
10
[ "PRISMA Cluster of Excellence and Mainz Institute for Theoretical Physics, Johannes Gutenberg-Universität Mainz, 55099 Mainz, Germany", "PRISMA Cluster of Excellence and Mainz Institute for Theoretical Physics, Johannes Gutenberg-Universität Mainz, 55099 Mainz, Germany", "PRISMA Cluster of Excellence and Mainz Institute for Theoretical Physics, Johannes Gutenberg-Universität Mainz, 55099 Mainz, Germany" ]
[ "2017JCAP...02..042H", "2017PhRvD..95g5001A", "2018JCAP...04..025B", "2019JCAP...11..011N", "2021MNRAS.503..458H", "2022JCAP...05..042G", "2022JHEP...10..186L", "2023PhRvD.107k5016N", "2024EPJC...84..136L", "2024arXiv240201839L" ]
[ "astronomy", "physics" ]
8
[ "High Energy Physics - Phenomenology", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1931AnP...403..257S", "1965TrAlm...5...87E", "1985ApJ...299.1001K", "1985PhRvL..55..257S", "1987ApJ...321..560G", "1987ApJ...321..571G", "1987NuPhB.279..804S", "1990ApJ...356..302G", "1996PhR...267..195J", "1998ASPC..131...76L", "2003JHEP...05..067B", "2003PhRvL..90v5002K", "2004PhRvD..69j1302B", "2005NatPh...1..147W", "2006AIPC..836..398W", "2006AJ....132.2685M", "2006PhRvL..96u1302F", "2007PhR...442..269W", "2008PhRvD..77a5008H", "2009JCAP...08..037N", "2009PhLB..674..250B", "2009PhRvD..80f3501Z", "2009PhRvD..80h3502K", "2010JCAP...01..018B", "2010JHEP...06..029M", "2010JPhG...37j5009C", "2010MNRAS.402...21N", "2010PhRvD..81f3502S", "2010PhRvD..81g5004B", "2010PhRvD..81h5032B", "2010PhRvD..82a5011M", "2010PhRvD..82h3509T", "2010PhRvD..82k5012S", "2011JCAP...04..003B", "2011JCAP...04..007E", "2011PhRvD..83f3521L", "2011PhRvD..84a4028F", "2012JCAP...05..002E", "2012JCAP...07..016A", "2012JCAP...10..045B", "2012PhRvD..86b3506S", "2012PhRvD..86d3527K", "2012PhRvD..86g6014F", "2013APh....43..171B", "2013APh....43..348F", "2013JCAP...07..010B", "2013JCAP...08..011B", "2013JCAP...11..032A", "2013PhRvD..87k5007T", "2013PhRvD..88e5005R", "2013PhRvD..88h3519H", "2014JCAP...03..053B", "2014JCAP...04..012I", "2014JCAP...10..049C", "2014PDU.....5...35D", "2014PhRvD..89d2001A", "2014PhRvD..90d3538B", "2014PhRvE..89c2145K", "2015JCAP...04..042C", "2015JCAP...05..047I", "2015JCAP...06..035K", "2015JCAP...08..036B", "2015JCAP...11..039R", "2015JHEP...08..050L", "2015NatPh..11..245I", "2015PNAS..11212249W", "2015PhRvD..91c5002B", "2015PhRvL.114e1301D", "2015PhRvL.114n1301C", "2016JCAP...04..022A", "2016JCAP...05..016A", "2016JCAP...11..007V", "2016PhRvD..93b3527S", "2016PhRvD..93e2014A", "2016PhRvD..93f1101A", "2016PhRvD..93k5036F", "2016PhRvD..94f3512M", "2016PhRvL.116p1301A", "2017PhLB..773..121A" ]
[ "10.1103/PhysRevD.95.055031", "10.48550/arXiv.1607.04278" ]
1607
1607.04278_arXiv.txt
16
7
1607.04278
Many theories of dark matter (DM) predict that DM particles can be captured by stars via scattering on ordinary matter. They subsequently condense into a DM core close to the center of the star and eventually annihilate. In this work, we trace DM capture and annihilation rates throughout the life of a massive star and show that this evolution culminates in an intense annihilation burst coincident with the death of the star in a core collapse supernova. The reason is that, along with the stellar interior, also its DM core heats up and contracts, so that the DM density increases rapidly during the final stages of stellar evolution. We argue that, counterintuitively, the annihilation burst is more intense if DM annihilation is a p -wave process than for s -wave annihilation because in the former case, more DM particles survive until the supernova. If among the DM annihilation products are particles like dark photons that can escape the exploding star and decay to standard model particles later, the annihilation burst results in a flash of gamma rays accompanying the supernova. For a galactic supernova, this "dark gamma-ray burst" may be observable in the Čerenkov Telescope Array.
false
[ "DM annihilation", "DM particles", "more DM particles", "DM", "DM capture", "s -wave annihilation", "annihilation rates", "gamma rays", "standard model particles", "dark matter", "stars", "ordinary matter", "particles", "stellar evolution", "the DM annihilation products", "dark photons", "a DM core", "its DM core", "an intense annihilation", "the DM density" ]
8.022881
-1.140205
-1
2878662
[ "Ho, I. -Ting", "Medling, Anne M.", "Groves, Brent", "Rich, Jeffrey A.", "Rupke, David S. N.", "Hampton, Elise", "Kewley, Lisa J.", "Bland-Hawthorn, Joss", "Croom, Scott M.", "Richards, Samuel", "Schaefer, Adam L.", "Sharp, Rob", "Sweet, Sarah M." ]
2016Ap&SS.361..280H
[ "LZIFU: an emission-line fitting toolkit for integral field spectroscopy data" ]
90
[ "Institute for Astronomy, University of Hawaii, Honolulu, HI, USA; Research School of Astronomy and Astrophysics, Australian National University, Weston, ACT, Australia", "Research School of Astronomy and Astrophysics, Australian National University, Weston, ACT, Australia", "Research School of Astronomy and Astrophysics, Australian National University, Weston, ACT, Australia", "Infrared Processing and Analysis Center, California Institute of Technology, Pasadena, CA, USA; Observatories of the Carnegie Institution of Washington, Pasadena, CA, USA", "Department of Physics, Rhodes College, Memphis, TN, USA", "Research School of Astronomy and Astrophysics, Australian National University, Weston, ACT, Australia", "Institute for Astronomy, University of Hawaii, Honolulu, HI, USA; Research School of Astronomy and Astrophysics, Australian National University, Weston, ACT, Australia", "Sydney Institute for Astronomy, School of Physics, University of Sydney, Sydney, NSW, Australia", "Sydney Institute for Astronomy, School of Physics, University of Sydney, Sydney, NSW, Australia; ARC Centre of Excellence for All-sky Astrophysics (CAASTRO), Sydney, NSW, Australia", "Sydney Institute for Astronomy, School of Physics, University of Sydney, Sydney, NSW, Australia; ARC Centre of Excellence for All-sky Astrophysics (CAASTRO), Sydney, NSW, Australia", "Sydney Institute for Astronomy, School of Physics, University of Sydney, Sydney, NSW, Australia; ARC Centre of Excellence for All-sky Astrophysics (CAASTRO), Sydney, NSW, Australia", "Research School of Astronomy and Astrophysics, Australian National University, Weston, ACT, Australia; ARC Centre of Excellence for All-sky Astrophysics (CAASTRO), Sydney, NSW, Australia", "Research School of Astronomy and Astrophysics, Australian National University, Weston, ACT, Australia" ]
[ "2016ApJ...827..103K", "2016MNRAS.462.1616D", "2016MNRAS.463..170C", "2017A&A...601A..61V", "2017ApJ...834..174K", "2017ApJ...846...39H", "2017ApJS..232...11T", "2017MNRAS.464..121S", "2017MNRAS.468.3965F", "2017MNRAS.470.1991B", "2017MNRAS.470.3395H", "2017MNRAS.470.4573Z", "2017MNRAS.470.4974D", "2017MNRAS.471.2438L", "2017MNRAS.472.2833S", "2018A&A...618A...6K", "2018A&A...618A..64H", "2018ApJ...852..113M", "2018ApJ...856...89T", "2018ApJ...863L..21K", "2018MNRAS.473..380T", "2018MNRAS.475..716G", "2018MNRAS.475.5194M", "2018MNRAS.479.4907D", "2018MNRAS.479.5235P", "2018MNRAS.481.2299S", "2019A&A...628A.117B", "2019AJ....158..160B", "2019ApJ...870...19G", "2019ApJ...873...52O", "2019ApJ...878..161J", "2019ApJ...885...32K", "2019ApJ...885L..31H", "2019ApJ...887...80K", "2019MNRAS.482.2166T", "2019MNRAS.483..458B", "2019MNRAS.483.2851S", "2019MNRAS.484.3042S", "2019MNRAS.485.4024V", "2019MNRAS.485L..38D", "2019MNRAS.487...79P", "2019MNRAS.487.4153D", "2019MNRAS.489.2830M", "2020A&A...634A.121H", "2020ARA&A..58...99S", "2020IAUS..352..144J", "2020MNRAS.495.3819A", "2020MNRAS.498.2372D", "2021A&A...645A.130C", "2021ApJ...911L..17R", "2021ApJ...923..160M", "2021MNRAS.502.3357P", "2021MNRAS.505..991C", "2021MNRAS.506.1295S", "2021MNRAS.508..489M", "2021RMxAA..57....3S", "2022A&A...659A.191E", "2022ApJ...925..203J", "2022ApJ...929..118G", "2022ApJ...933..110X", "2022ApJ...938...63L", "2022MNRAS.510..320F", "2022MNRAS.511.2105K", "2022MNRAS.512.1765O", "2022MNRAS.513.5988N", "2022MNRAS.514.4465M", "2022MNRAS.516.3411W", "2022MNRAS.516.3569B", "2022MNRAS.517.2677R", "2022NewA...9701895L", "2022arXiv221106005G", "2023A&A...670A...4B", "2023ApJ...954...77J", "2023IAUS..373..234G", "2023MNRAS.519.1098C", "2023MNRAS.519.1452W", "2023MNRAS.519.3237S", "2023MNRAS.519.4801C", "2023MNRAS.519.5324K", "2023MNRAS.521...84F", "2023MNRAS.524.1169L", "2023MNRAS.526.1613B", "2023NatAs...7..463R", "2024A&A...684A..53B", "2024ApJ...966..130R", "2024MNRAS.527.8566Z", "2024MNRAS.529..104M", "2024MNRAS.531.4017O", "2024arXiv240501637D", "2024arXiv240515072F" ]
[ "astronomy" ]
16
[ "Integral field spectroscopy", "Data analysis", "Emission line fitting", "Astrophysics - Astrophysics of Galaxies" ]
[ "1999lssu.conf..105C", "2000AJ....120.1579Y", "2001MNRAS.326...23B", "2002ApJ...571..545P", "2003MNRAS.342..345C", "2003SPIE.4834..161D", "2004PASP..116..138C", "2005MNRAS.357..945G", "2005MNRAS.358..363C", "2006MNRAS.365...46O", "2006MNRAS.366.1151S", "2006MNRAS.371..829S", "2006NewAR..49..501S", "2007Ap&SS.310..255D", "2007MNRAS.376..125S", "2007MNRAS.381.1252T", "2008MNRAS.385.1998K", "2009A&A...501.1269K", "2009ASPC..411..251M", "2009MNRAS.395...28M", "2010Ap&SS.327..245D", "2010ApJ...721..505R", "2010ApJ...723.1255R", "2010SPIE.7735E..08B", "2010SPIE.7735E..0MM", "2011AJ....141..192Z", "2011Ap&SS.331....1W", "2011ApJ...729L..27R", "2011MNRAS.413..813C", "2012A&A...538A...8S", "2012ApJ...752...97Y", "2012ApJ...753....5R", "2012MNRAS.419.1402G", "2012MNRAS.421..872C", "2012MNRAS.424..157V", "2012SPIE.8446E..0UG", "2012SPIE.8446E..13M", "2012SPIE.8446E..53L", "2013A&A...549A..87H", "2014A&A...561A.130C", "2014MNRAS.437.1070G", "2014MNRAS.444.3894H", "2014MNRAS.444.3961D", "2014MNRAS.445.1104R", "2015ApJ...798....7B", "2015ApJ...801...42D", "2015ApJS..217...12D", "2015IAUS..309...21B", "2015MNRAS.446.1567A", "2015MNRAS.446.2186M", "2015MNRAS.447.2857B", "2015MNRAS.448.2030H", "2015MNRAS.448.2301M", "2015MNRAS.450.2593V", "2015MNRAS.451.2780A", "2016MNRAS.457.1257H", "2016RMxAA..52...21S", "2016RMxAA..52..171S" ]
[ "10.1007/s10509-016-2865-2", "10.48550/arXiv.1607.06561" ]
1607
1607.06561_arXiv.txt
\label{ho16b-sec1} Galaxy emission-line spectroscopy has always been a powerful tool for the analysis of the dynamical, physical and chemical properties of galaxies. Traditionally, spectroscopy of galaxies has been obtained by dispersing the light either across a slit (sacrificing one spatial dimension) or from a fibre (producing a single integrated spectrum). Active development of modern integral field spectroscopy (IFS) has made capturing 3-dimensional structures of galaxies very efficient, revolutionising the way we observe and study galaxies. The complex and perhaps stochastic nature of different physical processes governing galaxy evolution has inspired large galaxy surveys. In recent decades, large fibre and slit spectroscopic surveys such as the Sloan Digital Sky Survey \citep[SDSS;][]{York:2000qe}, the 2dF Galaxy Redshift Survey \citep[][]{Colless:1999kl}, and the Deep Extragalactic Evolutionary Probe 2 survey \citep[DEEP2;][]{Davis:2003lr} have drastically improved our understanding of the global (unresolved) properties of galaxy populations at different epochs of the Universe. Integral field spectroscopic surveys have recently become feasible, providing access simultaneously to both spectral and kinematic information of large numbers of galaxies. Two pioneering IFS surveys, the SAURON survey \citep{Bacon:2001rt} and its extension the $\mathrm{ATLAS^{3D}}$ survey \citep{Cappellari:2011ys}, studied about 260 early type galaxies in the local Universe ($z<0.01$). Surveys targeting both the blue and red galaxy populations, such as the Calar Alto Legacy Integral Field Area (CALIFA) survey \citep{Sanchez:2012fj}, the Sydney-Australian-astronomical-observatory Multi-object Integral-field spectrograph (SAMI) Galaxy Survey \citep{Croom:2012qy,Bryant:2015bh}, and the Mapping Nearby Galaxies at Apache Point Observatory survey \citep[MaNGA;][]{Bundy:2015kx}, are currently underway. These IFS surveys will provide critical information to bridge the knowledge gaps resulting from the limited spatial and kinematic information delivered by previous single-fibre and slit spectroscopic surveys. The sample sizes and data flows of these modern IFS surveys are substantial. With each data cube containing typically one to two thousand spectra, the CALIFA survey plans to observe about 600 galaxies in the local Universe ($0.005<z<0.03$); the SAMI Galaxy Survey will reach a sample size of 3,400 galaxies at $z<0.12$; and the MaNGA survey will build up a sample of 10,000 galaxies at a similar redshift to the SAMI Galaxy Survey. Future surveys using high-multiplex integral field unit (IFU) instrument such as HECTOR on the Anglo-Australian Telescope will observe on the order of 100,000 galaxies \citep{Lawrence:2012fk,Bland-Hawthorn:2015dn}. Current and forthcoming wide-field IFU instruments are also delivering large quantity of high quality data, such as the Wide Field Spectrograph (WiFeS) on the Australian National University 2.3-m telescope \citep{Dopita:2007kx,Dopita:2010yq}, the new Multi Unit Spectroscopic Explorer (MUSE) on the Very Large Telescope \citep{Bacon:2010ph}, the SITELLE instrument on the Canada France Hawaii Telescope \citep{Grandmont:2012sp}, the Keck Cosmic Web Imager at the W. M. Keck Observatory \citep{Martin:2010th,Morrissey:2012ij}. Significant efforts have been placed in developing corresponding tools for analysing large volume of spectroscopic data. The stellar continuum contains valuable information about the stellar kinematics, chemistry and star formation history of galaxies. Packages such as the STEllar Content via Maximum A Posteriori \citep[{\scshape stecmap};][]{Ocvirk:2006fj} package, the penalized pixel-fitting \citep[\ppxf;][]{Cappellari:2004uq} routine and the {\scshape starlight} package \citep{Cid-Fernandes:2005fk} can perform spectral template fitting and extract various stellar properties. For investigating gas physics, the emission lines fitting tools such as the Gas AND Absorption Line Fitting code \citep[{\scshape gandalf};][]{Sarzi:2006lr}, the {\scshape fit3d} package (\citealt{Sanchez:2006uq,Sanchez:2007qy}; and the successor {\scshape pipe3d}; \citealt{Sanchez:2016fk,Sanchez:2016lr}), and the Peak ANalysis utility \citep[{\scshape pan}\footnote{{\scshape pan} was subsequently adapted and modified by Mark Westmoquette for astronomical requirements. See \url{http://ifs.wikidot.com/pan}. };][]{dimeo2005} are commonly adopted to measure emission line fluxes and kinematics. As the spectral resolution of the instruments continue to improve, the intrinsic non-Gaussian line profile complicates the emission line analysis. When the spectral resolution is high ($\rm R > 3000 $), galaxies with active gas dynamics, such as winds, outflows or AGN, usually present skewed line profiles that require fitting multiple, assumed Gaussian, components to separate the different kinematic components overlapping in the line-of-sight direction (also referred as ``spectral decomposition''). Performing spectral decomposition on large datasets is non-trivial as significant human input is usually required. Here, we present our emission line fitting pipeline LaZy-IFU\footnote{The framework of the \lzifu\ stems from {\scshape uhspecfit}, a tool developed at the University of Hawai\textquoteright i and employed in several previous spectroscopic studies on gas abundances and outflows \citep[e.g.,][]{Zahid:2011fj,Rupke:2010fk,Rich:2010yq,Rich:2012oq,Rupke:2011fk}.} (\lzifu; written in the Interactive Data Language [{\idl}]), which is designed to eliminate the need for individual treatment of each of many thousands of spectra across an IFS galaxy survey (such as CALIFA, SAMI or MaNGA). The main objective of \lzifu\ is to extract 2-dimensional emission line flux maps and kinematic maps useful for investigating gas physics in galaxies. \lzifu\ has already been adopted in various publications using data from multiple instruments and surveys, including MUSE \citep[][]{Kreckel:2016zl}, SAMI \citep[e.g.][]{Ho:2014uq,Richards:2014lr,Allen:2015dk,Ho:2016rf}, CALIFA \citep{Davies:2014fy,Ho:2015hl}, WiFeS \citep{Ho:2015hl,Dopita:2015lq,Dopita:2015eu,Vogt:2015ec,Medling:2015eu}, and SPIRAL on the Anglo-Australian Telescope \citep{McElroy:2015ve}. The following characteristics were considered carefully while developing \lzifu. First, the pipeline must perform spectral decomposition automatically without needing repeated human instructions. Second, the pipeline needs to be scriptable for batch reduction, such that when necessary the same results can be reproduced by re-executing the same scripts. Third, the pipeline must be flexible and generalised so that data from most modern IFS instruments can be accepted without major restructuring of the inputs. Finally, the calculation speed must be optimised and the pipeline has to take advantage of parallel processing because of 1) the huge data flow from multiplexed IFS surveys, and 2) the possibility of fitting the same datasets multiple times for various experimental purposes. The focus of this paper is to present the core structure of \lzifu\ (Section~\ref{ho16b-sec2}), and examine the errors produced by the pipeline (Section~\ref{ho16b-sec4}). We also show examples of applying \lzifu\ on the CALIFA survey and SAMI Galaxy survey (Section~\ref{ho16b-sec3}). Finally, the code will be continuously maintained and made available to the public through github (\url{https://github.com/hoiting/LZIFU/releases}). We discuss future plans for the code in Section~\ref{ho16b-sec5}.
\label{ho16b-sec5} We have presented \lzifu, an \idl\ toolkit for fitting multiple emission lines and constructing emission line flux maps and kinematic maps from IFS data. We outlined the structure of \lzifu, and described in detail how the code performs spectral fitting and decomposition. We have also conducted simulations to examine the errors estimated by \lzifu\ and discussed the its limitations. We have demonstrated how \lzifu\ can be adopted to analyse data from the CALIFA survey and the SAMI Galaxy Survey. In some applications, single component fitting is adequate to capture the dominant kinematic component (typically from \ion{H}{ii} regions tracing disk rotation) and can produce flux and kinematic maps useful for various studies of gas physics. In cases where the line profiles are more sophisticated due to either active physical environments (e.g., AGN or galactic wind) or observational effects (e.g., beam smearing), multiple component fitting can better constrain the total line fluxes and provide more insight into the various physical processes. Although only examples from CALIFA and SAMI were presented in the paper, \lzifu\ is by no mean limited to these two datasets. Data from world-class IFS instruments with distinct structures (i.e. fibre-based, image-splitting), sizes and spatial resolutions have already been processed by \lzifu, including MUSE, WiFeS and SPIRAL. The \lzifu\ products and scientific results extracted from these can be found in \citet{Ho:2015hl,Dopita:2015eu,McElroy:2015ve,Vogt:2015ec,Kreckel:2016zl}. Further data from these instruments, and surveys from other IFS instruments are currently being analysed, with a wealth of scientific results from \lzifu\ products expected to be published in the coming years. While this paper outlines the official release version of \lzifu, future improvements to the pipeline will be implemented. For example, in a soon-to-be available upgrade we will include the option of fitting binned data. Spatially binning data can significantly improve the detection of faint emission lines at large galactic radii. Different binning schemes such as contour binning \citep{Sanders:2006lr} and Voronoi tessellations \citep{Cappellari:2003yq} have established the usefulness of binning imaging and IFS data. On longer timescales, we plan to incorporate a full Bayesian analysis such that the parameter space can be explored more thoroughly and more accurate errors can be reported. It is also possible to analyse mock 3D data cubes from numerical simulations parallel to observational data cubes to directly compare similar parameter maps. \acknowledgment We thank the referee for constructive comments that improve the quality of this work. LJK gratefully acknowledges the support of an ARC Future Fellowship, and ARC Discovery Project DP130103925. SMC acknowledges the support of an Australian Research Council Future Fellowship (FT100100457). The SAMI Galaxy Survey is based on observations made at the Anglo-Australian Telescope. The Sydney-AAO Multi-object Integral field spectrograph was developed jointly by the University of Sydney and the Australian Astronomical Observatory. The SAMI input catalogue is based on data taken from the Sloan Digital Sky Survey, the GAMA Survey and the VST ATLAS Survey. The SAMI Galaxy Survey is funded by the Australian Research Council Centre of Excellence for All-sky Astrophysics, through project number CE110001020, and other participating institutions. The SAMI Galaxy Survey website is \url{http://sami-survey.org/}. This study makes uses of the data provided by the Calar Alto Legacy Integral Field Area (CALIFA) survey (\url{http://califa.caha.es/}). Based on observations collected at the Centro Astron\'{o}mico Hispano Alem\'{a}n (CAHA) at Calar Alto, operated jointly by the Max-Planck-Institut f\"{u}r Astronomie and the Instituto de Astrofisica de Andalucia (CSIC). \appendix \twocolumn
16
7
1607.06561
We present lzifu (LaZy-IFU), an idl toolkit for fitting multiple emission lines simultaneously in integral field spectroscopy (IFS) data. lzifu is useful for the investigation of the dynamical, physical and chemical properties of gas in galaxies. lzifu has already been applied to many world-class IFS instruments and large IFS surveys, including the Wide Field Spectrograph, the new Multi Unit Spectroscopic Explorer (MUSE), the Calar Alto Legacy Integral Field Area (CALIFA) survey, the Sydney-Australian-astronomical-observatory Multi-object Integral-field spectrograph (SAMI) Galaxy Survey. Here we describe in detail the structure of the toolkit, and how the line fluxes and flux uncertainties are determined, including the possibility of having multiple distinct kinematic components. We quantify the performance of lzifu, demonstrating its accuracy and robustness. We also show examples of applying lzifu to CALIFA and SAMI data to construct emission line and kinematic maps, and investigate complex, skewed line profiles presented in IFS data. The code is made available to the astronomy community through github. lzifu will be further developed over time to other IFS instruments, and to provide even more accurate line and uncertainty estimates.
false
[ "emission line", "fitting multiple emission lines", "IFS data", "large IFS surveys", "multiple distinct kinematic components", "other IFS instruments", "integral field", "IFS", "kinematic maps", "Multi Unit Spectroscopic Explorer", "flux uncertainties", "complex, skewed line profiles", "SAMI", "the Calar Alto Legacy Integral Field Area", "CALIFA and SAMI data", "the line fluxes", "many world-class IFS instruments", "the new Multi Unit Spectroscopic Explorer", "galaxies", "the Calar Alto Legacy Integral Field Area (CALIFA) survey" ]
11.914273
6.994238
196
12516930
[ "Huby, Elsa", "Absil, Olivier", "Mawet, Dimitri", "Baudoz, Pierre", "Femenıa Castellã, Bruno", "Bottom, Michael", "Ngo, Henry", "Serabyn, Eugene" ]
2016SPIE.9909E..20H
[ "The QACITS pointing sensor: from theory to on-sky operation on Keck/NIRC2" ]
4
[ "Univ. de Liège (Belgium)", "Univ. de Liège (Belgium)", "California Institute of Technology (United States)", "LESIA, Observatoire de Paris, CNRS, UPMC, Univ. Paris-Diderot (France)", "W. M. Keck Observatory (United States)", "California Institute of Technology (United States)", "California Institute of Technology (United States)", "Jet Propulsion Lab. (United States)" ]
[ "2016SPIE.9909E..22F", "2017AJ....153...43S", "2022JATIS...8b9006V", "2023arXiv231206806B" ]
[ "astronomy", "physics" ]
7
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "2005ApJ...633.1191M", "2010ApJ...709...53M", "2012A&A...539A.126M", "2012SPIE.8442E..04M", "2012SPIE.8446E..8KD", "2013A&A...552L..13M", "2014SPIE.9151E..19F", "2015A&A...584A..74H", "2015PASP..127..890J" ]
[ "10.1117/12.2233274", "10.48550/arXiv.1607.05497" ]
1607
1607.05497_arXiv.txt
\label{sec:intro} Small inner working angle (IWA) coronagraphs are the key to access the full angular resolution potential of current large ground based telescopes. The IWA is defined as the angular separation at which the flux of an off-axis companion is transmitted by 50\%. Among the existing solutions, the vector vortex coronagraph is based on a focal plane phase mask inducing a phase ramp onto the beam. This kind of coronagraph is attractive for several reasons, including high extinction ratio, achromatic properties, a continuous discovery space and small IWA. For all these reasons, this kind of coronagraph can be found in several leading instruments: the Palomar infrared camera PHARO \cite{Mawet2010a}, Subaru/SCExAO \cite{Jovanovic2015}, VLT/VISIR \cite{Delacroix2012,Kerber2014}, VLT/NACO\cite{Mawet2013}, LBT/LMIRCam \cite{Defrere2014} and recently Keck/NIRC2 (first light obtained in June 2015, see Ref.~\citenum{FemeniaCastella2016} in these proceedings and Serabyn et al. in prep). The latter four are mid-infrared instruments working in the L or N band and are equipped with Annular Groove Phase Masks\cite{Mawet2005} (AGPM). These components are developed by the University of Li\`ege and manufactured by the University of Uppsala\cite{Forsberg2014}. For a review of the results obtained with these instruments, see Ref.~\citenum{Absil2016} in these proceedings. However, a small IWA focal-plane coronagraph is also a synonym for high sensitivity to pointing error. Indeed, a slight shift of the star from the mask center will result in a starlight leakage and degrade the contrast performance. A low order wavefront sensor is therefore required in high contrast imaging instruments (for a full review of low order wavefront sensor possibilities, see Ref.~\citenum{Mawet2012}), in addition to the Adaptive Optics system. Fig.~\ref{fig:transmission} shows the experimental transmission measured in L band with an AGPM. This curve is valid for an off-axis companion as well as for the central star. In order to estimate the non common path aberrations, this additional sensor has to be placed as close as possible to the coronagraphic mask and/or science camera. In the case of the vortex phase mask, we have developed a pointing sensor algorithm called QACITS\cite{Huby2015} (Quadrant Analysis of Coronagraphic Images for Tip-tilt Sensing). While the principle has first been empirically introduced and implemented in laboratory for the Four Quadrant Phase Mask\cite{Mas2012}, we have derived the complete theoretical framework adapted to the vortex coronagraph\cite{Huby2015} and successfully implemented it on-sky. \begin{figure} \begin{center} \begin{tabular}{c} \includegraphics[height=6cm]{vortex_mask_exp-transmission-lin} \end{tabular} \end{center} \caption[transmission] { \label{fig:transmission} Experimental transmission of a vortex phase mask (AGPM of topological charge 2) for a circular non obstructed pupil. The flux is integrated on a disk of diameter equal to the full width at half maximum. Estimated IWA is 0.9\,$\lambda/D$ as predicted by simulations.} \end{figure}
The QACITS algorithm has been specifically developed for stabilizing the positioning of a star onto the center of a vortex coronagraph. This algorithm is necessary to optimize the extinction performance of the coronagraph. In addition, the procedure that we have implemented at Keck/NIRC2 is fully automated and takes care of every steps of the acquisition: calibration, centering optimization, science frames acquisition. As a result, the vortex observation mode is now "user-friendly" and optimized in efficiency (manual alignment can be much longer and tiresome for the observer). Besides, the ease of operation allowed by the QACITS procedure has made possible to offer the vortex mode for science in shared risk mode since 2016B. The simplicity of its implementation and its robustness make QACITS very attractive for other existing vortex modes. Preliminary tests have been carried out at LBT/LMIRCam (in collaboration with D.\,Defr\`ere from Universit\'e de Li\`ege), and at VLT/NACO (in collaboration with J.\,Girard and G.\,Zins from ESO). The results are very promising, and the implementation of a NACO template dedicated to observation with the vortex mode is currently under consideration.
16
7
1607.05497
Small inner working angle coronagraphs are essential to benefit from the full potential of large and future extremely large ground-based telescopes, especially in the context of the detection and characterization of exoplanets. Among existing solutions, the vortex coronagraph stands as one of the most effective and promising solutions. However, for focal-plane coronagraph, a small inner working angle comes necessarily at the cost of a high sensitivity to pointing errors. This is the reason why a pointing control system is imperative to stabilize the star on the vortex center against pointing drifts due to mechanical flexures, that generally occur during observation due for instance to temperature and/or gravity variations. We have therefore developed a technique called QACITS<SUP>1</SUP> (Quadrant Analysis of Coronagraphic Images for Tip-tilt Sensing), which is based on the analysis of the coronagraphic image shape to infer the amount of pointing error. It has been shown that the flux gradient in the image is directly related to the amount of tip-tilt affecting the beam. The main advantage of this technique is that it does not require any additional setup and can thus be easily implemented on all current facilities equipped with a vortex phase mask. In this paper, we focus on the implementation of the QACITS sensor at Keck/NIRC2, where an L-band AGPM has been recently commissioned (June and October 2015), successfully validating the QACITS estimator in the case of a centrally obstructed pupil. The algorithm has been designed to be easily handled by any user observing in vortex mode, which is available for science in shared risk mode since 2016B.
false
[ "shared risk mode", "vortex mode", "science", "Small inner working angle coronagraphs", "exoplanets", "mechanical flexures", "existing solutions", "error", "errors", "Coronagraphic Images", "QACITS", "a vortex phase mask", "temperature and/or gravity variations", "characterization", "instance", "observation", "the vortex coronagraph", "October", "a pointing control system", "June" ]
14.388288
10.692083
10
12608612
[ "Boncioli, Denise", "Fedynitch, Anatoli", "Winter, Walter" ]
2017NatSR...7.4882B
[ "Nuclear Physics Meets the Sources of the Ultra-High Energy Cosmic Rays" ]
33
[ "DESY, Platanenallee 6, D-15738, Zeuthen, Germany", "DESY, Platanenallee 6, D-15738, Zeuthen, Germany", "DESY, Platanenallee 6, D-15738, Zeuthen, Germany" ]
[ "2017EPJWC.14107001B", "2017JCAP...01..033B", "2017JCAP...11..009A", "2017PhRvD..95l3001L", "2018A&A...611A.101B", "2018APh...102...39H", "2018ApJ...854...54R", "2018MNRAS.476.1191B", "2018NPNew..28...12R", "2018NatSR...810828B", "2018PhRvD..98d3001S", "2019ApJ...872..110B", "2019ApJ...873...88H", "2019EPJWC.20804002B", "2019EPJWC.20808001P", "2019JCAP...01..002A", "2019JCAP...05..006A", "2019JCAP...05..047M", "2019JCAP...11..007M", "2019PhR...801....1A", "2020MNRAS.498.5990H", "2020PhR...872....1B", "2020PhRvD.102l3008B", "2022JHEP...06..105V", "2022NatRP...4..697G", "2022arXiv221003756B", "2023APh...15202866K", "2023ApJ...948...42W", "2023JCAP...05..024A", "2023PPNL...20..637K", "2023PhRvD.107d3019V", "2023PhRvD.107h3008F", "2023arXiv231200409A" ]
[ "astronomy", "physics", "general" ]
4
[ "Astrophysics - High Energy Astrophysical Phenomena", "High Energy Physics - Phenomenology", "Nuclear Experiment", "Nuclear Theory" ]
[ "1948PhRv...74.1046G", "1976ApJ...205..638P", "1996PhDT........59R", "2000CoPhC.124..290M", "2005AIPC..769..199K", "2005AIPC..769.1303F", "2005APh....23..191K", "2006RPPh...69.2259M", "2008APh....29....1A", "2008JCAP...10..033A", "2008PhRvD..78b3005M", "2010ApJ...721..630H", "2010PhRvD..81l3001M", "2011PhRvD..84j5007T", "2012A&A...539A..88C", "2012JCAP...10..007A", "2012PhRvL.108w1101H", "2013APh....42...41K", "2013ApJ...768..186B", "2013JCAP...03..010F", "2014APh....54...48T", "2014JCAP...10..020A", "2014NDS...120..211B", "2014NDS...120..272O", "2014PhRvD..90l2005A", "2014PhRvD..90l2006A", "2015APh....62...66B", "2015EPJA...51..185F", "2015JCAP...07..042T", "2015JCAP...09..023G", "2015JCAP...10..034K", "2015JCAP...10..063A", "2015MNRAS.451..751G", "2015NatCo...6.6783B", "2015PhRvD..92f3011T", "2015PhRvD..92l3001U", "2016JCAP...05..038A", "2016PhLB..762..288A", "2017JCAP...04..038A", "2018A&A...611A.101B" ]
[ "10.1038/s41598-017-05120-7", "10.48550/arXiv.1607.07989" ]
1607
1607.07989_arXiv.txt
Particles from space reaching the Earth with energies higher than $10^{9}$ GeV are detected by ultra-high energy cosmic ray (UHECR) observatories such as the Pierre Auger Observatory~\cite{ThePierreAuger:2015rma} and the Telescope Array (TA) experiment~\cite{AbuZayyad:2012kk}. UHECRs are expected to be accelerated in astrophysical sources and to travel through extragalactic space before they hit the Earth's atmosphere; they can interact with photons in both environments. The primary composition of UHECRs is still unknown; however, the mass composition measured by the Auger Observatory indicates heavier elements at the highest energies beyond $10^{9.3}$~GeV~\cite{Aab:2014kda,Aab:2014aea,Aab:2016htd,Porcelli:2015}, \ie, significantly heavier than helium and at most as heavy as iron. The study of interactions of nuclei is therefore critical for our understanding of cosmic ray astrophysics both within sources and during propagation. Most of the literature, as for example \cite{Allard:2008gj,Taylor:2011ta,Taylor:2013gga,Taylor:2015rla,Fang:2013cba,Aloisio:2013hya,diMatteo:2015,Peixoto:2015ava,Aab:2016zth}, focuses on finding the right cosmic ray composition injected from the sources into the intergalactic medium, propagating it through the cosmic microwave background (CMB) and the extragalactic background light (EBL), which are thermal target photon fields, \ie, relatively strongly peaked. The long-term vision is, however, to trace back the cosmic ray composition into the \revise{sources}, which requires a combined source--propagation model; see \eg\ \Refs~\cite{Globus:2014fka,Unger:2015laa}. Such models face several challenges, including 1) largely unknown astrophysical environments and uncertainties, 2) limited computational resources for detailed simulations and parameter space studies, and 3) a complicated interplay of radiation processes -- especially in the source, where the photon fields are often non-thermal (power laws) and the physical processes have been less studied. One of the main challenges is therefore to walk the line between precision and efficiency to overcome problems 2) and 3), while new insights on the astrophysical parameters 1) are to be obtained from multi-messenger observations. An example is \Ref~\cite{Baerwald:2014zga}, where, for a pure proton composition, constraints on the astrophysical parameters are derived from cosmic ray and neutrino observations. \revise{In order to extend such approaches to heavier compositions}, the physical processes have to be controlled with a precision as high as possible, where the photo-disintegration of nuclei plays the leading role. While the required target precision is arguable, a first bottleneck is the description of the source as accurate as the propagation from source to detector -- where the target photon environment can be very different, and which has been well studied. In this work, we focus on the photo-disintegration of nuclei, which has been extensively studied in the CMB and EBL, where it is the dominant process changing the mass composition of the nuclei. The leading contribution to photo-disintegration is an excitation called ``giant dipole resonance'' (GDR)~\cite{Goldhaber:1948zza}, which can be interpreted as a vibration of the bulk of protons and neutrons leading to a resonant structure. This process occurs above $\sim$8~MeV (energy in the nucleus' rest frame) and causes the disruption of the primary nucleus with the emission of one or two nucleons. At higher energies the ``quasi-deuteron'' (QD) process dominates, where the photon interacts with a nucleon pair followed by consequent ejection of nucleons or light fragments. Note that we do not consider astrophysical situations where disintegration is dominated by even higher energy processes, such as baryonic resonances, at energies beyond 150~MeV. A frequently used model in cosmic ray astrophysics is Puget-Stecker-Bredekamp (PSB)~\cite{Puget:1976nz}, that relies on choosing one isotope for each mass number $A$, and a unique disintegration chain populated through subsequent emission of nucleons. This approach is implemented for cosmic ray propagation in the {\it SimProp} software~\cite{Aloisio:2012wj}. A more sophisticated approach, based on the TALYS nuclear reaction program~\cite{Koning:2007}, is implemented in the cosmic ray propagation software CRPropa2 and 3~\cite{Kampert:2012fi,Batista:2016yrx}, which includes 183 isotopes and 2200 channels for the photo-disintegration. Differences in modeling the interactions affect the observables (as energy spectrum and composition), and consequently have an impact on the interpretation of UHECR measurements~\cite{Aab:2016zth,diMatteo:2015,Batista:2015mea}. UHECRs are expected to be accelerated in astrophysical sources, such as Gamma-Ray Bursts \revise{(GRBs; see \Ref~\cite{Meszaros:2006rc} for a review), Active Galactic Nuclei (AGNs), starburst galaxies, or jets produced in other cataclysmic events (mergers of neutron stars or black holes, tidal disruptions of massive stars getting too close to a super-massive black hole, \etc) -- to name a few examples. In sources such as GRBs or AGNs, they will disintegrate in the strong photon field present in the jets.} Examples for disintegration treatments \revise{in the sources} include~\cite{Globus:2014fka,Anchordoqui:2007tn}, where the GDR modeling follows \cite{Khan:2004nd} with a modified Gaussian parametrization of cross sections from \cite{Puget:1976nz}. The GDR resonance is even more simplified in \cite{Murase:2008mr,Murase:2010gj,Bustamante:2014oka} as a box function. Other authors use semi-analytical implementations of existing UHECR propagation codes (such as CRPropa)~\cite{Unger:2015laa}. So far, however, the astrophysical sources have not been simulated with a complexity comparable to CRPropa for cosmic ray propagation -- including several hundred isotopes and the ten thousands of disintegration channels among them. That can be attributed to the fact that the target photon spectrum, relevant for the photo-disintegration, is {\em a priori} arbitrary, \ie, it can have a very different shape compared to that of the CMB. In this work, we present a description of the processes in the sources with a level of complexity of the interactions comparable to that of the most sophisticated cosmic ray propagation models. We review the available nuclear data necessary to construct and verify reliable interaction models, and we point out what information is missing from nuclear physics. \begin{figure*}[t!] \begin{center} \includegraphics[width=0.7\textwidth]{exfor_chart.pdf} \end{center} \caption{\label{fig:exfor_chart} Experimental situation versus astrophysical requirements for nuclear isotopes interesting for cosmic ray astrophysics (gray boxes, from TALYS \cite{Koning:2007} and CRPropa2 \cite{Kampert:2012fi}). Experimental measurements (from EXFOR database \cite{Otuka2014272}) are marked by red (yellow) boxes if the total absorption (any inclusive cross section) has been measured. Theoretical models are marked by dots (ENDF \cite{Chadwick20112887}, JENDL \cite{JENDL-PD-2004}, PEANUT \cite{FLUKA_pd1,FLUKA_pd2}, PSB \cite{Puget:1976nz}). Calculations in cosmic ray astrophysics require the total and inclusive cross sections for the blue isotopes. \revise{These isotopes have been obtained by recursively following all possible paths from all possible injection elements with different threshold multiplicities (for priorities~1 and~2), see main text for details.} Injected isotopes are framed by black rectangles (\revise{we inject the most abundant stable isotope for each $Z$}). \revise{Violet framed rectangles} refer to very unstable isotopes integrated out in the disintegration chain.} \end{figure*}
Since the evidence for a heavy composition of the observed cosmic rays is condensing, understanding the origin of cosmic rays requires the description of astrophysical sources with high radiation densities in the presence of nuclei. The long-term vision to determine the injected cosmic ray composition \revise{in the sources} and to identify unknown astrophysical parameters by multi-messenger astronomy requires combined source-propagation models for cosmic ray nuclei. As a first step into that direction, we have identified the requirements from nuclear physics for the leading process governing the nuclear cascade both within the sources and during cosmic ray propagation -- which correspond to very different (non-thermal versus thermal) radiation fields. Our working hypothesis has been that both sources and propagation need to be described by a comparable level of complexity for these applications, and we have therefore developed \revise{new methods} capabable to describe the radiation processes in the sources including the full nuclear disintegration chain \revise{efficiently}. As one key result, we have compared the situation of nuclear measurements (red/yellow in \figu{exfor_chart}) with the input needed for cosmic ray astrophysics (blue in \figu{exfor_chart}). We have demonstrated that the measurements on the participating nuclear isotopes are sparse, as unstable nuclei are produced in the photo-disintegration -- which can live relatively long at extremely high energies where their lifetimes are Lorentz-boosted. Although sophisticated nuclear models exist, such as TALYS, their prediction power for the considered isotopes seems limited to within about a factor of two. While the impact on individual photo-disintegration rates can be large, we have found that random fluctuations tend to average out in the nuclear cascade unless there are systematic effects not probably accounted for, such as offsets between neutron- and proton-rich elements, along the main diagonal, or a missed feeddown into lighter isotopes. We therefore propose systematic measurements to improve the predictability of unmeasured cross sections. A long term goal could be measurements of total and inclusive cross sections of the blue-marked isotopes in \figu{exfor_chart}, or corresponding predictive information from nuclear theory. Finally, we have demonstrated that -- while simple nuclear disintegration models for the sources may be sufficient in some cases -- there are observables which require a treatment at the level of complexity of TALYS or FLUKA within the cosmic ray source. For example, the ejected cosmic ray composition from a GRB deviates in the energy range where intermediate mass nuclei are produced in these models up to a factor of two in $\langle \ln A \rangle$. Similar consequences are expected for other observables, such as the secondary neutrino production, which requires further study. We conclude that nuclear cosmic ray astrophysics will emerge as a new discipline which faces a new level of complexity compared to proton-only radiation models for the sources. If we really want to understand the physics of the sources and close the argumentation chain from the cosmic ray injection \revise{within the source}, over particle acceleration and radiation, to cosmic ray propagation, we need better models especially on the source side. We have presented a radiation model with an unprecedented level of complexity for the nuclear disintegration in a gamma-ray burst, where the developed methods can be used in combined source-propagation models in the future, or in studies of the particle acceleration itself. While the astrophysical parameters and environments are highly uncertain, one can use such models to address the reverse question, \ie, use multi-messenger astronomy to derive the injection composition in the sources and to study the source parameters. \revise{The radiation density in most of the possible UHECR sources dominates over the matter density; for this reason, the photonuclear reactions are the responsible for energy losses and production of secondary particles. However, for the case of Galactic cosmic rays, interactions with the interstellar medium have to be considered during propagation in the Galaxy. A similarity with what we studied in the present work }can be found from the point of view of the importance of nuclear physics. Secondary nuclei and antiparticles are produced by spallation of primary cosmic rays and carry information about origin of the primaries and transport in the Galaxy \cite{Coste:2011jc}. In particular, the most relevant uncertainties in modeling the antiproton flux, that constitutes also one of the prime channels for indirect searches of Dark Matter, come from the absence of measurements of p-He reactions \cite{Giesen:2015ufa,Kappl:2015bqa}. As a consequence, new input from nuclear physics could be of crucial importance also in \revise{nuclei-nuclei/nucleon processes, mostly relevant for} Galactic cosmic ray astrophysics.\\
16
7
1607.07989
The determination of the injection composition of cosmic ray nuclei within astrophysical sources requires sufficiently accurate descriptions of the source physics and the propagation - apart from controlling astrophysical uncertainties. We therefore study the implications of nuclear data and models for cosmic ray astrophysics, which involves the photo-disintegration of nuclei up to iron in astrophysical environments. We demonstrate that the impact of nuclear model uncertainties is potentially larger in environments with non-thermal radiation fields than in the cosmic microwave background. We also study the impact of nuclear models on the nuclear cascade in a gamma-ray burst radiation field, simulated at a level of complexity comparable to the most precise cosmic ray propagation code. We conclude with an isotope chart describing which information is in principle necessary to describe nuclear interactions in cosmic ray sources and propagation.
false
[ "cosmic ray sources", "cosmic ray nuclei", "astrophysical sources", "non-thermal radiation fields", "astrophysical environments", "cosmic ray astrophysics", "astrophysical uncertainties", "nuclear model uncertainties", "propagation", "environments", "nuclear models", "the most precise cosmic ray propagation code", "nuclear interactions", "the cosmic microwave background", "nuclear data", "nuclei", "a gamma-ray burst radiation field", "complexity", "the source physics", "models" ]
6.072927
0.401148
13
12510090
[ "Pavluchenko, Sergey A." ]
2016PhRvD..94h4019P
[ "Cosmological dynamics of spatially flat Einstein-Gauss-Bonnet models in various dimensions: Low-dimensional Λ -term case" ]
25
[ "Programa de Pós-Graduação em Física, Universidade Federal do Maranhão (UFMA), 65085-580 São Luís, Maranhão, Brazil" ]
[ "2017EPJC...77...89E", "2017EPJC...77..402E", "2017EPJC...77..503P", "2017GrCo...23..359C", "2017MPLA...3250202E", "2018EPJC...78..100I", "2018EPJC...78..373P", "2018EPJC...78..551P", "2018EPJC...78..611P", "2018EPJWC.16802003T", "2018GReGr..50..119I", "2018GrCo...24...28C", "2018Parti...1....4P", "2019EPJC...79..111P", "2019GrCo...25..164E", "2019MPLA...3450111E", "2020AnPhy.41968216C", "2020EPJC...80..543E", "2020GrCo...26...16I", "2020Symm...12..250E", "2021EPJC...81..136C", "2021arXiv210410423P", "2022RSPTA.38010177I", "2022Symm...14.1296E", "2023Symm...15..783I" ]
[ "astronomy", "physics" ]
4
[ "High Energy Physics - Theory", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
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[ "10.1103/PhysRevD.94.084019", "10.48550/arXiv.1607.07347" ]
1607
1607.07347_arXiv.txt
Einstein's general relativity was formulated more than hundred years ago, but the extra-dimensional models are even older. Indeed, the first attempt to construct an extra-dimensional model was performed by Nordstr\"om~\cite{Nord1914} in 1914. It was a vector theory that unified Nordstr\"om's second gravity theory~\cite{Nord_2grav} with Maxwell's electromagnetism. Later in 1915 Einstein introduced General Relativity (GR)~\cite{einst}, but still it took almost four years to prove that Nordstr\"om's theory and others were wrong. During the solar eclipse of 1919, the bending of light near the Sun was measured and the deflection angle was in perfect agreement with GR, while Nordstr\"om's theory, being scalar gravity, predicted a zeroth deflection angle. But Nordstr\"om's idea about extra dimensions survived, and in 1919 Kaluza proposed~\cite{KK1} a similar model but based on GR: in his model five-dimensional Einstein equations could be decomposed into four-dimensional Einstein equations plus Maxwell's electromagnetism. In order to perform such a decomposition, the extra dimensions should be ``curled'' or compactified into a circle and ``cylindrical conditions'' should be imposed. Later in 1926, Klein proposed~\cite{KK2, KK3} a nice quantum mechanical interpretation of this extra dimension and so the theory called Kaluza-Klein was formally formulated. Back then their theory unified all known interactions at that time. With time, more interactions were known and it became clear that to unify them all, more extra dimensions are needed. Nowadays, one of the promising theories to unify all interactions is M/string theory. Presence in the Lagrangian of the curvature-squared corrections is one of the distinguishing features of the string theories gravitational counterpart. Indeed, Scherk and Schwarz~\cite{sch-sch} were the first to discover the potential presence of the $R^2$ and $R_{\mu \nu} R^{\mu \nu}$ terms in the Lagrangian of the Virasoro-Shapiro model~\cite{VSh1, VSh2}. A curvature squared term of the $R^{\mu \nu \lambda \rho} R_{\mu \nu \lambda \rho}$ type appears~\cite{Candelas_etal} in the low-energy limit of the $E_8 \times E_8$ heterotic superstring theory~\cite{Gross_etal} to match the kinetic term for the Yang-Mills field. Later it was demonstrated~\cite{zwiebach} that the only combination of quadratic terms that leads to a ghost-free nontrivial gravitation interaction is the Gauss-Bonnet (GB) term: $$ L_{GB} = L_2 = R_{\mu \nu \lambda \rho} R^{\mu \nu \lambda \rho} - 4 R_{\mu \nu} R^{\mu \nu} + R^2. $$ \noindent This term, first found by Lanczos~\cite{Lanczos1, Lanczos2} (therefore it is sometimes referred to as the Lanczos term) is an Euler topological invariant in (3+1)-dimensional space-time, but not in (4+1) and higher dimensions. Zumino~\cite{zumino} extended Zwiebach's result on higher-than-squared curvature terms, supporting the idea that the low-energy limit of the unified theory might have a Lagrangian density as a sum of contributions of different powers of curvature. In this regard the Einstein-Gauss-Bonnet (EGB) gravity could be seen as a subcase of more general Lovelock gravity~\cite{Lovelock}, but in the current paper we restrain ourselves with only quadratic corrections and so to the EGB case. Extra-dimensional theories have one thing in common---one needs to explain where additional dimensions are ``hiding'', since we do not sense them, at least with the current level of experiments. One of the possible ways to hide extra dimensions, as well as to recover four-dimensional physics, is to build a so-called ``spontaneous compactification'' solution. Exact static solutions with the metric being a cross product of a (3+1)-dimensional manifold and a constant curvature ``inner space'' were build for the first time in~\cite{add_1}, but with the (3+1)-dimensional manifold being Minkowski (the generalization for a constant curvature Lorentzian manifold was done in~\cite{Deruelle2}). In the context of cosmology, it is more interesting to consider a spontaneous compactification in the case where the four-dimensional part is given by a Friedmann-Robertson-Walker metric. In this case it is also natural to consider the size of the extra dimensions as time dependent rather than static. Indeed in \cite{add_4} it was explicitly shown that in order to have a more realistic model one needs to consider the dynamical evolution of the extra-dimensional scale factor. In~\cite{Deruelle2}, the equations of motion for compactification with both time-dependent scale factors were written for arbitrary Lovelock order in the special case of a spatially flat metric (the results were further proven in~\cite{prd09}). The results of~\cite{Deruelle2} were reanalyzed for the special case of 10 space-time dimensions in~\cite{add_10}. In~\cite{add_8}, the existence of dynamical compactification solutions was studied with the use of Hamiltonian formalism. More recently, efforts to find spontaneous compactifications were made in~\cite{add13}, where the dynamical compactification of the (5+1) Einstein-Gauss-Bonnet model was considered; in \cite{MO04, MO14} with different metric {\it Ans\"atze} for scale factors corresponding to (3+1)- and extra-dimensional parts; and in \cite{CGP1, CGP2, CGPT}, where general (e.g., without any {\it Ansatz}) scale factors and curved manifolds were considered. Also, apart from cosmology, the recent analysis has focused on properties of black holes in Gauss-Bonnet~\cite{alpha_12, add_rec_1, add_rec_2, addn_1, addn_2} and Lovelock~\cite{add_rec_3, add_rec_4, addn_3, addn_4} gravities, features of gravitational collapse in these theories~\cite{addn_5, addn_6, addn_7}, general features of spherical-symmetric solutions~\cite{addn_8}, and many others. In the context of finding exact solutions, the most common {\it Ansatz} used for the functional form of the scale factor is exponential or power law. Exact solutions with exponential functions for both the (3+1)- and extra-dimensional scale factors were studied for the first time in \cite{Is86}, and an exponentially increasing (3+1)-dimensional scale factor and an exponentially shrinking extra-dimensional scale factor were described. Power-law solutions have been analyzed in \cite{Deruelle1, Deruelle2} and more recently in~\cite{mpla09, prd09, Ivashchuk, prd10, grg10} so that there is an almost complete description (see also~\cite{PT} for useful comments regarding physical branches of the solutions). Solutions with exponential scale factors~\cite{KPT} have been studied in detail, namely, models with both variable~\cite{CPT1} and constant~\cite{CST2} volume, developing a general scheme for constructing solutions in EGB; recently~\cite{CPT3} this scheme was generalized for general Lovelock gravity of any order and in any dimensions. Also, the stability of the solutions was addressed in~\cite{my15} (see also~\cite{iv16} for stability of general exponential solutions in EGB gravity), where it was demonstrated that only a handful of the solutions could be called ``stable'', while the remaining are either unstable or have neutral/marginal stability, and so additional investigation is required. In order to find all possible regimes of Einstein-Gauss-Bonnet cosmology, it is necessary to go beyond an exponential or power-law {\it Ansatz} and keep the functional form of the scale factor generic. We are also particularly interested in models that allow dynamical compactification, so it is natural to consider the metric as the product of a spatially three-dimensional part and an extra-dimensional part. In that case the three-dimensional part represents ``our Universe'' and we expect for this part to expand while the extra-dimensional part should be suppressed in size with respect to the three-dimensional one. In \cite{CGP1} it was found that there exists a phenomenologically sensible regime in the case when the curvature of the extra dimensions is negative and the Einstein-Gauss-Bonnet theory does not admit a maximally symmetric solution. In this case the three-dimensional Hubble parameter and the extra-dimensional scale factor asymptotically tend to the constant values. In \cite{CGP2} a detailed analysis of the cosmological dynamics in this model with generic couplings was performed. Recently this model was also studied in~\cite{CGPT}, where it was demonstrated that, with an additional constraint on couplings, Friedmann-type late-time behavior could be restored. The current paper is a spiritual successor of~\cite{my16a}, where we investigated cosmological dynamics of the vacuum Einstein-Gauss-Bonnet model. In both papers the spatial section is a product of two spatially flat manifolds with one of them three-dimensional, which represents our Universe and the other is extra-dimensional. In~\cite{my16a} we considered vacuum model while in the current paper -- the model with the cosmological term. In~\cite{my16a} we demonstrated that the vacuum model has two physically viable regimes -- first of them is the smooth transition from high-energy GB Kasner to low-energy GR Kasner. This regime appears for $\alpha > 0$ at $D=1,\,2$ and for $\alpha < 0$ at $D \geqslant 2$ (so that at $D=2$ it appears for both signs of $\alpha$). The other viable regime is smooth transition from high-energy GB Kasner to anisotropic exponential regime with expanding three-dimensional section (``our Universe'') and contracting extra dimensions; this regime occurs only for $\alpha > 0$ and at $D \geqslant 2$. Let us note that in~\cite{CGP1, CGP2, CGPT} we considered similar model but with both manifolds to be constant (generally non-zero) curvature. Unlike the paper with vacuum solutions, in this paper we limit ourselves with only lower-dimensional ($D=1$ and $D=2$) cases; the higher-dimensional cases -- $D=3$ and the general $D\geqslant 4$ case -- will be considered in a separate forthcoming paper. The structure of the manuscript is as follows: first we write down general equations of motion for Einstein-Gauss-Bonnet gravity, then we rewrite them for our symmetry {\it Ansatz}. In the following sections we analyze them for $D=1$ and $D=2$ cases, considering the $\Lambda$-term case in this paper only. Each case is followed by a small discussion of the results and properties of this particular case; after considering all cases we discuss their properties, generalities, and differences and draw conclusions.
To conclude, we performed thorough analysis of $\Lambda$-term regimes in Einstein-Gauss-Bonnet gravity in two lowest number of dimensions -- five and six. We have considered the manifold which is a product of three-dimensional (which represents our Universe) and extra-dimensional (in our case with $D=1,\,2$ dimensions) parts. This separation is quite natural as with it we could describe natural compactification. Our analysis demonstrate that generally $\Lambda$-term models have much mode abundant dynamics then vacuum cases~\cite{my16a}. Our investigation also suggest that in $D=1$ model there are no physically viable regimes. On contrary, $D=2$ case have smooth transitions from high-energy Kasner regime to anisotropic exponential solutions with contracting extra dimensions. In one particular case $\alpha\Lambda = -3/2$ with $\alpha < 0$ and $\Lambda > 0$ the size of extra dimensions (in the sense of the scale factor) reaches constant value (and so the expansion rate $h(t)\to 0$), making this case similar to the regime described in spatially-curved ``geometric frustration'' model~\cite{CGP1}. Both $D=1$ and $D=2$ cases lack regular low-energy regime -- in the former of them it is singular (one faces finite-time singularity while reaching low-energy Kasner regime) and for the latter $H=0$ is not an endpoint and the evolution continues to $H<0$ domain until either nonstandard singularity or exponential solution is reached. So that in $D=2$ case we have interesting regimes like the transition from isotropic exponential contracting to isotropic exponential expansion (like a bounce) and anisotropic regimes with contracting three-dimensional spaces turn to expansion and vice versa. Lack of low-energy regimes in $D=2$ as well as presence of $h=0$ anisotropic exponential solution have the same cause -- $h(H)$ expression in $D=2$ case is distinct from both $D=1$ and the entire vacuum case~\cite{my16a}. Indeed, in both $D=1$ and the vacuum case, $h(H)$ and $H(t)$ curves have either $h\to 0$ or $h\to\pm\infty$ as $H\to 0$. In the former case we have low-energy nonsingular Kasner regime, in the latter -- the same but singular. Also one can see that we cannot have $h=0$ while $H\ne 0$, which prevent corresponding exponential solutions to exist. On contrary, one can see that in $D=2$ we have $h\ne 0$ at $H=0$ so that $h=0$, $H\ne 0$ exponential solutions exist while low-energy regimes are absent. With the same argumentation, $h=0$, $H\ne 0$ exponential solutions could exist in higher-order Lovelock models as well, except for lowest possible dimensions, like $D=3$ for cubic Lovelock, $D=4$ for quadric and so on. Overall, present study brought us several unexpected results -- for power-law solutions, we found singular low-energy Kasner-like behavior for $D=1$ and Milne-like behavior for $D=2$. Both regimes are supposed to be forbidden to exist in presence of Lambda-term, yet the analysis in term of Kasner exponents points on them. Of course, none of these regimes are reached, but the fact that analysis points on them could indicate that they formally could exist but are unstable -- so this situation is in need for the additional investigation. Exponential solutions also behave not exactly as expected -- multiple distinct isotropic solutions for both $D=1$ and $D=2$ as well as directional stability of anisotropic solutions in $D=2$ clearly indicate need of additional study of exponential solutions as well. Low-dimensional $\Lambda$-term case demonstrated interesting dynamics for both $D=1$ and $D=2$ with some unexpected features. In forthcoming paper we consider $D=3$ and generic $D \geqslant 4$ $\Lambda$-term cases and finalize our study $\Lambda$-term case.
16
7
1607.07347
In this paper we perform a systematic study of spatially flat [(3 +D )+1 ] -dimensional Einstein-Gauss-Bonnet cosmological models with the Λ -term. We consider models that topologically are the product of two flat isotropic subspaces with different scale factors. One of these subspaces is three-dimensional and represents our space and the other is D -dimensional and represents extra dimensions. We consider no Ansatz on the scale factors, which makes our results quite general. With both Einstein-Hilbert and Gauss-Bonnet contributions in play, the cases with D =1 and D =2 have different dynamics due to the different structure of the equations of motion. We analytically study equations of motion in both cases and describe all possible regimes. It is demonstrated that the D =1 case does not have physically viable regimes while D =2 has a smooth transition from high-energy Kasner to an anisotropic exponential regime. This transition occurs for two ranges of α and Λ : α &gt;0 with α Λ ≤1 /2 (including Λ &lt;0 ) and α &lt;0 , Λ &gt;0 with α Λ &lt;-3 /2 . For the latter case, if α Λ =-3 /2 , the extra-dimensional part has h →0 and so the size of extra dimensions (in the sense of the scale factor) is reaching a constant value. We report substantial differences between D =1 and D =2 cases and between these cases and their vacuum counterparts, describe features of the cases under study, and discuss the origin of the differences.
false
[ "D =", "different scale factors", "extra dimensions", "Λ ≤1", "Λ", "different dynamics", "lt;0", "=", "an anisotropic exponential regime", "substantial differences", "the extra-dimensional part", "motion", "all possible regimes", "physically viable regimes", "study", "Λ &lt;0", "equations", "Kasner", "the scale factor", "the scale factors" ]
10.272448
0.25372
89
12434660
[ "Hoeneisen, B." ]
2016arXiv160702424H
[ "Study of baryon acoustic oscillations with SDSS DR12 data and measurement of $\\Omega_\\textrm{DE}(a)$" ]
2
[ "-" ]
[ "2016arXiv160808486H", "2017IJAA....7...11H" ]
[ "astronomy" ]
4
[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
[ "2000astro.ph..9071H", "2007ApJ...664..660E", "2008PhRvD..77d3525S", "2010deot.book..246B", "2013PhR...530...87W", "2015ApJS..219...12A" ]
[ "10.48550/arXiv.1607.02424" ]
1607
1607.02424_arXiv.txt
A point-like peak in the primordial density of the universe results, well after recombination and decoupling, in two spherical shells of overdensity: one of radius $\approx 150$ Mpc, and one of radius $\approx 18$ Mpc \cite{Eisenstein, BAO1, BAO2}. (All distances in this article are ``co-moving", i.e. are referred to the present time $t_0$.) The ``acoustic length scale" $r'_S \approx 150$ Mpc is approximately the distance that sound waves of the tightly coupled plasma of photons, electrons, protons, and helium nuclei traveled from the time of the Big Bang until the electrons recombined with the protons and helium nuclei to form neutral atoms, and the photons decoupled. The inner spherical shell of $\approx 18$ Mpc becomes re-processed by the hierarchical formation of galaxies \cite{BH}, while the radius of the outer shell is unprocessed to better than 1\% \cite{BAO2} (or even 0.1\% with corrections \cite{BAO2}) and therefore is an excellent standard ruler to measure the expansion of the universe as a function of redshift $z$. Histograms of galaxy-galaxy distances show an excess of galaxy pairs with distances in the approximate range $150 - 18$ to $150 + 18$ Mpc. This ``Baryon Acoustic Oscillation" (BAO) signal has a signal-to-background ratio of the order of 0.1\% to 1\%. Measurements of these BAO signals are by now well established (see \cite{BAO1, BAO2} for extensive lists of early publications). In this article we present studies of BAO with Sloan Digital Sky Survey (SDSS) data release DR12 \cite{DR12}.
The main results of these studies are the 18 independent BAO distance measurements presented in Table \ref{d0} which do not depend on any cosmological parameter. These BAO distance measurements alone place a strong constraint on $\Omega_\textrm{DE} + 0.5 \Omega_k = 0.646 \pm 0.022$ (for constant $\Omega_\textrm{DE}$) or $0.656 \pm 0.065$ (when $\Omega_\textrm{DE}(a)$ is allowed to depend on $a$), while the constraint on $\Omega_k$ is weak. Constraints on the cosmological parameters in several scenarios from the BAO measurements alone and from BAO plus $\theta_\textrm{MC}$ measurements are presented in Tables \ref{BAO_fit}, \ref{BAO_thetaMC_fit_12} and \ref{BAO_thetaMC_fit_00}. We find that the 18 BAO distance measurements plus $\theta_\textrm{DE}$ are in agreement with a family of universes with different $\Omega_k$. Two examples are $\Omega_k = 0.065$ with $\Omega_\textrm{DE}(a)$ constant, and $\Omega_k = 0$ with $\Omega_\textrm{DE}(a)$ varying significantly with $a$. For these two examples we present measurements of $\Omega_\textrm{DE}(a)$ in Figs. \ref{O_DE_z} and \ref{O_DE_z_all}. The cosmology with both $\Omega_k = 0$ and $\Omega_\textrm{DE}(a)$ constant has a tension of $4.3 \sigma$ with the BAO plus $\theta_\textrm{MC}$ data. The BAO plus $\theta_\textrm{MC}$ data for constant $\Omega_\textrm{DE}(a)$ obtains $\Omega_k = 0.061 \pm 0.012$ for $\Delta d_z = 0.0012$ and $\Omega_k = 0.037 \pm 0.011$ for $\Delta d_z = 0$. These results have some tension with independent observations: $\Omega_k = -0.042^{+0.024}_{-0.022}$ from CMB data alone with the assumption of constant $\Omega_\textrm{DE}(a)$ \cite{PDG}, and $\Omega_k \approx 0.000 \pm 0.013$ from CMB plus supernova (SN) data with the assumption of constant $\Omega_\textrm{DE}(a)$ \cite{PDG}. If $\Omega_\textrm{DE}(a)$ is allowed to be variable there are only loose constraints on $\Omega_k$ from independent measurements, so the case with $\Omega_k = 0$ with non-constant $\Omega_\textrm{DE}(a)$ is viable.
16
7
1607.02424
We define Baryon Acoustic Oscillation (BAO) distances $\hat{d}_\alpha(z, z_c)$, $\hat{d}_z(z, z_c)$, and $\hat{d}_/(z, z_c)$ that do not depend on cosmological parameters. These BAO distances are measured as a function of redshift $z$ with the Sloan Digital Sky Survey (SDSS) data release DR12. From these BAO distances alone, or together with the correlation angle $\theta_\textrm{MC}$ of the Cosmic Microwave Background (CMB), we constrain the cosmological parameters in several scenarios. We find $4.3 \sigma$ tension between the BAO plus $\theta_\textrm{MC}$ data and a cosmology with flat space and constant dark energy density $\Omega_\textrm{DE}(a)$. Releasing one and/or the other of these constraints obtains agreement with the data. We measure $\Omega_\textrm{DE}(a)$ as a function of $a$.
false
[ "several scenarios", "constant dark energy density", "cosmological parameters", "flat space", "DR12", "agreement", "z_c)$", "\\Omega_\\textrm{DE}(a)$.", "SDSS", "Baryon Acoustic Oscillation", "the Sloan Digital Sky Survey", "(SDSS) data release", "the cosmological parameters", "BAO", "CMB", "the data", "\\hat{d}_/(z", "\\hat{d}_\\alpha(z", "redshift", "the Cosmic Microwave Background" ]
12.127378
2.288707
161
12613181
[ "Le Tiec, Alexandre", "Novak, Jérôme" ]
2017ogw..book....1L
[ "Theory of Gravitational Waves" ]
5
[ "LUTH, Observatoire de Paris, PSL Research University, CNRS, Université Paris Diderot, Sorbonne Paris Cité, 5 place Jules Janssen, 92195 Meudon Cedex, France", "LUTH, Observatoire de Paris, PSL Research University, CNRS, Université Paris Diderot, Sorbonne Paris Cité, 5 place Jules Janssen, 92195 Meudon Cedex, France" ]
[ "2016PhRvD..94h4018A", "2018AnPhy.394..225V", "2018arXiv181004477B", "2023arXiv230900772A", "2024EPJC...84..406D" ]
[ "astronomy", "physics" ]
24
[ "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1972gcpa.book.....W", "1973grav.book.....M", "1984ucp..book.....W", "2003gieg.book.....H", "2004sgig.book.....C", "2007tste.book.....K", "2008LRR....11....8L", "2008PhRvL.100d1101S", "2009fcgr.book.....S", "2011ApJ...743..102G", "2011gwpa.book.....C", "2013RvMP...85.1401A", "2013gere.book.....S", "2013srgf.book.....G", "2014LRR....17....2B", "2015CQGra..32g4001L", "2016LRR....19....1A" ]
[ "10.1142/9789813141766_0001", "10.48550/arXiv.1607.04202" ]
1607
1607.04202_arXiv.txt
\label{s:th_intro} Together with black holes and the expansion of the Universe, the existence of gravitational radiation is one of the key predictions of Einstein's general theory of relativity.\cite{Ei.16,Ei.18} The discovery of the binary pulsar PSR B1913+16,\cite{HuTa.75} and the subsequent observation of its orbital decay, as well as that of other binary pulsars, have provided strong evidence for the existence of gravitational waves.\cite{WeHu.16,Lo.08} These observations have triggered an ongoing international effort to detect gravitational waves directly, mainly by using kilometer-scale laser interferometric antennas such as the LIGO and Virgo detectors.\cite{Aa.al.15,Ac.al.15} During the months of September and October 2015, the Advanced LIGO antennas have detected, for the first time, gravitational waves generated by two distinct cosmic sources. These waves were emitted, more than a billion years ago, during the coalescence of two binary black hole systems of $65M_\odot$ and $22M_\odot$, respectively.\cite{Ab.al2.16,Ab.al3.16} Many more gravitational-wave observations are expected to follow before the end of this decade.\cite{Ab.al.16} These are truly exciting times, because the direct observation of gravitational waves is going to have a tremendous impact on physics, astrophysics and cosmology.\cite{SaSc.09} In this chapter, we provide a short but self-contained introduction to the theory of gravitational waves. No prior knowledge of general relativity shall be assumed, and only those concepts that are necessary for an introductory discussion of gravitational radiation will be introduced. For more extensive treatments, the reader is referred to the resource letter \refcite{Ce.03}, the review articles \refcite{Th.87,ScRi.01,FlHu.05,Bu.07,Bl.14,BuSa.15}, and the topical books \refcite{Mag,CrAn}. Most general relativity textbooks include a discussion of gravitational radiation, such as Refs.~\refcite{Wei,MTW,Wal,Har,Car,Schu,Str}. The remainder of this chapter is organized as follows. Section \ref{s:intro_gw} provides a qualitative introduction to gravitational waves. Section \ref{s:geo} introduces the geometrical setting (manifold, metric, connection) that is required to formulate the general theory of relativity, the topic of Sec.~\ref{s:Einstein}. Then, gravitational waves are defined, in Sec.~\ref{s:def_gw}, as solutions of the linearized Einstein equation around flat (Minkowski) spacetime. These waves are shown to propagate at the speed of light and to possess two polarization states. The interaction of gravitational waves with matter, an important topic that underlies their direct detection, is addressed in Sec.~\ref{s:th_inter}. Finally, Sec.~\ref{s:gen} provides an overview of the generation of gravitational radiation by matter sources. In particular, Einstein's quadrupole formulas are used to show, using order-of-magnitude estimates, that nonspherical compact objects moving at relativistic speeds are powerful gravitational wave emitters. Throughout this chapter we use units in which $c = 1$, except in Secs.~\ref{s:intro_gw} and \ref{s:gen}, where we keep all occurences of the speed of light. Our conventions are those of Ref.~\refcite{MTW}; in particular, we use a metric signature $-,+,+,+$.
16
7
1607.04202
The existence of gravitational radiation is a natural prediction of any relativistic description of the gravitational interaction. In this chapter, we focus on gravitational waves, as predicted by Einstein's general theory of relativity. First, we introduce those mathematical concepts that are necessary to properly formulate the physical theory, such as the notions of manifold, vector, tensor, metric, connection and curvature. Second, we motivate, formulate and then discuss Einstein's equation, which relates the geometry of spacetime to its matter content. Gravitational waves are later introduced as solutions of the linearized Einstein equation around flat spacetime. These waves are shown to propagate at the speed of light and to possess two polarization states. Gravitational waves can interact with matter, allowing for their direct detection by means of laser interferometers. Finally, Einstein's quadrupole formulas are derived and used to show that nonspherical compact objects moving at relativistic speeds are powerful gravitational wave sources.
false
[ "laser interferometers", "powerful gravitational wave sources", "Gravitational waves", "gravitational waves", "relativistic speeds", "curvature", "gravitational radiation", "flat spacetime", "nonspherical compact objects", "matter", "spacetime", "connection", "Einstein", "tensor", "vector", "means", "relativity", "Einsteins general theory", "its matter content", "the linearized Einstein equation" ]
9.355021
1.414392
70
746093
[ "Mauerhan, Jon", "Williams, G. Grant", "Smith, Nathan", "Smith, Paul S.", "Filippenko, Alexei V.", "Hoffman, Jennifer L.", "Milne, Peter", "Leonard, Douglas C.", "Clubb, Kelsey I.", "Fox, Ori D.", "Kelly, Patrick L." ]
2014MNRAS.442.1166M
[ "Multi-epoch spectropolarimetry of SN 2009ip: direct evidence for aspherical circumstellar material" ]
91
[ "Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA; Steward Observatory, University of Arizona, 933 N. Cherry Ave., Tucson, AZ 85721, USA", "Steward Observatory, University of Arizona, 933 N. Cherry Ave., Tucson, AZ 85721, USA; MMT Observatory, Tucson, AZ 85721-0065, USA", "Steward Observatory, University of Arizona, 933 N. Cherry Ave., Tucson, AZ 85721, USA", "Steward Observatory, University of Arizona, 933 N. Cherry Ave., Tucson, AZ 85721, USA", "Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA", "Department of Physics and Astronomy, University of Denver, 2112 East Wesley Avenue, Denver, CO 80208, USA", "Steward Observatory, University of Arizona, 933 N. Cherry Ave., Tucson, AZ 85721, USA", "Department of Astronomy, San Diego State University, PA-210, 5500 Campanile Drive, San Diego, CA 92182-1221, USA", "Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA", "Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA", "Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA" ]
[ "2014A&A...571A..89Q", "2014ApJ...789..104O", "2014MNRAS.441.2230H", "2015A&A...578A.100W", "2015AJ....149....9M", "2015ApJ...803L..26M", "2015EAS....71..163H", "2015IAUS..305..269H", "2015MNRAS.447..598S", "2015MNRAS.447..772F", "2015MNRAS.447.1922M", "2015MNRAS.449.1876S", "2015MNRAS.450..246B", "2015MNRAS.453.3886F", "2015MNRAS.453.4467M", "2015PASA...32...19A", "2015arXiv151100784H", "2016ApJ...817...66K", "2016AstBu..71..422G", "2016MNRAS.457.3241A", "2016MNRAS.459.1039T", "2016MNRAS.463.2904S", "2016MNRAS.463.3894E", "2016PhDT.......149T", "2016arXiv160900731G", "2017A&A...599A.129T", "2017ApJ...834..118M", "2017ApJ...835..140M", "2017IAUS..329...54H", "2017MNRAS.469.1559G", "2017MNRAS.470.1491R", "2017MNRAS.471.4047A", "2017RSPTA.37560268S", "2017RSPTA.37560273T", "2017arXiv170108885R", "2017hsn..book..403S", "2017hsn..book..693V", "2017hsn..book..967W", "2017suex.book.....B", "2018ApJ...853...57B", "2018ApJ...853L...8B", "2018ApJ...856...29M", "2018MNRAS.473.4805K", "2018MNRAS.474..197P", "2018MNRAS.474.4988S", "2018MNRAS.475.1104B", "2018MNRAS.477...74A", "2018MNRAS.479.2249S", "2018SPIE10704E..2HW", "2019A&A...622A..93B", "2019ApJ...872..135K", "2019ApJ...872..141S", "2019ApJ...874...80M", "2019ApJ...882...70M", "2019ApJ...885...43A", "2019ApJ...887..249S", "2019BAAA...61...90Q", "2019BAAS...51c.339G", "2019MNRAS.482.4057M", "2019MNRAS.482.4233G", "2019MNRAS.488.3089K", "2019MNRAS.490.4536Q", "2020ApJ...892...13S", "2020ApJ...899...51S", "2020MNRAS.492.5897S", "2020MNRAS.497.5395N", "2020MNRAS.498.3835B", "2020MNRAS.499.3544S", "2020RSOS....700467F", "2021ApJ...921L..35L", "2021MNRAS.502..176C", "2021MNRAS.503..312M", "2021MNRAS.505.3664N", "2021MNRAS.507.1229P", "2022A&A...662L..10R", "2022ApJ...928..138P", "2022ApJ...935L..33J", "2022ApJ...936...28T", "2022ApJ...936..114M", "2022ApJ...938...84D", "2022AstL...48..442C", "2022ChPhC..46c0006L", "2022MNRAS.513.5666B", "2022MNRAS.515...71S", "2023A&A...677A.105D", "2023A&A...677L...1P", "2023ApJ...955L..37V", "2024ApJ...964..181H", "2024MNRAS.527.6090M", "2024MNRAS.529.1104B", "2024MNRAS.530..405S" ]
[ "astronomy" ]
8
[ "circumstellar matter", "stars: evolution", "supernovae: general", "supernovae: individual: SN 2009ip", "stars: winds", "outflows", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1093/mnras/stu730", "10.48550/arXiv.1403.4240" ]
1403
1403.4240_arXiv.txt
Interacting supernovae (SNe) are stellar explosions that collide with dense circumstellar material (CSM) produced by the progenitor star. These events are raising critical new questions about the final evolutionary phases of massive stars and the mass-loss episodes that ensue before core collapse (Smith \& Arnett 2014). Type-IIn and Ibn SNe, in particular, are interacting SNe that are characterized spectroscopically by the presence of relatively narrow emission lines of H and He in their spectra (Schlegel 1990; Filippenko 1997; Pastorello et al. 2008), which arise from dense CSM that becomes illuminated by the shock between the fast moving SN ejecta and slower moving CSM (Chevalier \& Fransson 1994). As such, observations of interacting SNe probe the stellar progenitor's pre-SN mass-loss history, providing valuable information on its final evolutionary episodes. Various lines of evidence show that interacting SNe require eruptive pre-SN mass loss that is reminiscent of luminous blue variable (LBV) stars, like $\eta$ Car, although observations indicate a wide range of mass-loss properties (e.g., see the review by Smith 2014). The eruptions are often detectable as extragalactic transients, commonly referred to as `SN impostors' (Van Dyk 2000; Smith et al. 2011a; Kochanek et al. 2012), which can rival the luminosity of a SN and spectroscopically mimic SNe~IIn. Although the observational distinction between interacting SNe and SN impostors is not always clear, a direct link connecting these phenomena has been established for several objects. The progenitors of SN~2006jc and SN~2011ht were detected undergoing luminous outbursts 1--2~yr prior to their core-collapse explosions (Foley et al. 2007; Pastorello et al. 2008; Fraser et al. 2013a). SN~2010mc (Ofek et al. 2013a) was also detected 1--2 months before dramatically brightening, although for this case it appears that the precursor event may actually have been the initial phases of the SN caught unusually early (Smith, Mauerhan, \& Prieto 2014). Perhaps the most interesting transient observed to have undergone multiple phases of eruptive mass loss is SN~2009ip. Originally classified as a SN (Maza et al. 2009), this object was actually discovered during an LBV outburst --- that is, as a SN~impostor (Smith et al. 2010; Foley et al. 2011). Several years of continued activity were followed photometrically and spectroscopically, leading up to SN~2009ip's most extreme outburst in 2012 (Mauerhan et al. 2013; Prieto et al. 2013; Levesque et al. 2014; Smith et al. 2013b; Smith, Mauerhan, \& Prieto 2014; Pastorello et al. 2013; Fraser et al. 2013b; Margutti et al. 2014; Graham et al. 2014). The 2012 event was comprised of two main components: the ``2012a" phase, marked by an initial peak at $M\approx-15.5$ mag that lasted for just over 1 month; and the ``2012b" phase, which began with a fast 10-day rise to a second peak of $M=-18.5$ mag, followed by a bumpy decline down to a slowly declining floor in the light curve near $M\approx-13$ mag. Based on spectral similarities with known core-collapse SNe, Mauerhan et al. (2013a) suggested that the relatively faint 2012a phase of SN~2009ip marked the initial stages of a core-collapse SN, while the subsequent 2012b brightening was the result of interaction between this SN and dense CSM ejected during the earlier LBV outbursts. Levesque et al. (2014) also shared the view that the 2012b brightening was the result of interaction between the 2012a outflow and existing CSM. However, several authors have suggested potential alternatives to a core collapse explosion -- that SN~2009ip's 2012 evolution was possibly the result of one or more nonterminal outbursts (Pastorello et al. 2013; Fraser et al. 2013b; Margutti et al. 2014), perhaps caused by the pulsational pair-instability mechanism. The motivation for nonterminal scenarios has been based largely on the fact that the total radiated energy of the light curve ($1.3\times10^{49}$ erg) can be explained by an explosion energy of $<10^{51}$~ergs (if spherical symmetry is assumed for the CSM), and also because the late-time data were interpreted as looking different from what is expected for radioactive decay phases. More recently, Smith, Mauerhan \& Prieto (2014; hereafter SMP14) showed that the 2012 light curve and spectral evolution are consistent with published models for core-collapse SNe, which can be initially faint (like SN~1987A) owing to a relatively compact progenitor radius (i.e., a blue supergiant, as expected for an LBV, instead of a red supergiant), while the late-time characteristics could be explained if the mass of synthesized $^{56}$Ni was half that of SN~1987A. SMP14 further argued that the radiated energy did not provide an argument against core collapse, since the CSM is probably aspherical, leading to an inefficient conversion of kinetic energy into radiation. In line with the spectral modeling results and interpretation of Levesque et al. (2014), which are consistent with a toroidal distribution of CSM, SMP14 also proposed a disk-like distribution of CSM, but further argued that significant asphericity is required by the fact that broad photospheric lines are still seen at late times, even after strong CSM interaction ends. This requires that a large fraction of the total solid angle of ejecta was able to expand unimpeded by CSM. SMP14 also pointed out that the $\sim$100-day persistence of broad photospheric lines requires a large ejecta mass of at least a few M$_{\odot}$, and is incompatible with a 0.1~M$_{\odot}$ shell, which would become optically thin much more quickly. A mass of a few M$_{\odot}$ moving at high speeds ($\sim$8000~km~s$^{-1}$) directly implies $\gtrsim10^{51}$~ergs of kinetic energy. In the case of SNe~IIn, the system geometry is critical if the total radiated energy from CSM interaction is to be used to infer the kinetic energy of the SN explosion. Fortunately, spectropolarimetry can yield important clues about the geometry of SNe, allowing us to stringently test the hypothesis of aspherical CSM. Here, we report multi-epoch spectropolarimetry of the 2012 evolution of SN~2009ip, from the end of the initial 2012a phase, through the peak of 2012b, and later into its decline. The results unambiguously demonstrate that the source of bright continuum emission during the 2012b phase (i.e., the CSM interaction zone) was aspherical, consistent with a toroidal or disk-like distribution of CSM proposed by SMP14 and Levesque et al. (2014). Furthermore, the results indicate that the initial 2012a phase has a separate polarization component having an axis of symmetry that is distinct from 2012b, which implies that the 2012a event did not create the CSM responsible for the 2012b brightening. Diminishing polarization at late times, in addition to similarities in wavelength-dependent polarization across lines for SN~2009ip and other SNe, provides additional evidence for a persistent underlying SN photosphere. Altogether, the available evidence supports the hypothesis of a SN explosion for SN~2009ip in 2012, and argues strongly against nonterminal explosion models for this object (at least those proposed thus far).
We have presented multi-epoch spectropolarimetry of the 2012 outburst of SN~2009ip, covering the 2012a and 2012b phases. Since the available data imply a very low amount of interstellar absorption and ISP $<0.2$\%, we have not attempted to remove ISP in our analysis. This could potentially lead to some small systematic errors in our results, although we expect any such changes will be minor and will not significantly influence our interpretation. The degree of polarization at the peak of 2012b implies an approximate axial ratio of $<$0.7 (probably $<$0.6) for the polarized CSM, which is consistent with the toroidal distribution of material suggested previously (SMP14; Levesque et al. 2014). In the future, modeling of the polarization properties of shocked toroidal CSM could reveal how well the comparison to the oblate electron-scattering atmosphere models of H{\"o}filch (1991) can characterise the axial ratio of such geometry. Details aside, the spectropolarimetric results demand substantial asphericity for the source of the 2012b luminosity and are thus inconsistent with models for SN~2009ip's 2012 explosion that invoke a spherical distribution of CSM (Fraser et al. 2013; Margutti et al. 2014). For the 2012a phase, the detection of significant polarization at a position angle that is nearly orthogonal to 2012b is consistent with a bipolar geometry for the SN outflow or partial obscuration of a potentially more spherical SN photosphere by the toroidal CSM. In any case, the orthogonality is inconsistent with the launching and collision of successive shells of CSM (Pastorello et al. 2012) generated $\approx$40 days apart (Margutti et al. 2014), since such a scenario should produce similar axes of symmetry for the 2012a and 2012b phases, not orthogonal geometries. Furthermore, the geometric parameters constrained by spectropolarimetry imply that only a small fraction of the SN ejecta ($<$10\%) participated in strong interaction with the toroidal CSM. Therefore, earlier estimates of a $\sim10^{50}$~ergs explosion that are based on the total radiated energy, which assume spherical symmetry for the CSM (Fraser et al. 2013; Margutti et al. 2014), are likely to be substantially underestimated by an order of magnitude. The strong similarities between SN~2009ip and SN~1993J, from the time near their peak through their decline, provide another strong line of evidence for the existence of an underlying SN photosphere from high-velocity ejecta (5000--8000~km~s$^{-1}$) that persists throughout SN~2009ip's 2012 evolution ($\sim$100 days). Such a long-lasting optically thick component indicates a mass of at least a few M$_{\odot}$ for the fast outer component of the ejecta, and the associated speeds imply a kinetic energy of $\gtrsim10^{51}$~ergs for the ejecta, as suggested by SMP14. After peak polarization, higher-order structure in the $Q$--$U$ plane from He-{\sc i}/Na~D indicates the presence of fast metal-rich ejecta that have overrun the dense CSM, consistent with a high-velocity flow that bypassed intense CSM interaction. Meanwhile, the weakening of polarization with declining continuum flux at late times suggests that the diminishing intensity of electron scattering is responsible for the decline in polarization. The fact that the flux and polarization drop more rapidly for SN~2009ip relative to other SNe~IIn, such as SN~1997eg (Hoffman et al. 2008) or SN~2010jl (Grant Williams, 2014, private communication), also suggests that the CSM around SN~2009ip has a compact configuration, compared to these other explosions. The lines of evidence presented here thus favour a scenario first proposed by Mauerhan et al. (2013), in which the luminosity of the 2012a phase of SN~2009ip was the result of a $\gtrsim10^{51}$~ergs explosion (i.e., a core-collapse SN) that subsequently plowed into an aspherical (probably toroidal) distribution of CSM during 2012b, creating the jump in luminosity and the strong polarization of the source. Nonterminal scenarios are difficult to support in light of the current body of evidence, and we would have to invent a process capable of generating ejecta having $\gtrsim10^{51}$~ergs of kinetic energy without destroying the star. Still, the best evidence of progenitor death is a vanishing of the star at late times. However, CSM interaction can persist for decades or more, and generate enough luminosity to make progenitor disappearance a difficult observable to confirm. SN~1961V is a particularly controversial example (Filippenko et al. 1995; Chu et al. 2004; Smith et al. 2011b; Kochanek et al. 2011; Van Dyk et al. 2012). Moreover, it is very plausible that the progenitor of SN~2009ip's 2012 explosion left a massive companion at the explosion site, not only because massive stars are rarely solitary (Sana et al. 2012), but because toroidal or disk-like geometry for the CSM is suggestive of binary influence. As discussed by Smith~\&~Arnett (2014), repeated brief pre-SN eruptions could result from binary interaction if the primary suddenly increases its radius during late nuclear burning stages, thereby triggering binary interaction and mass ejection. Finally, we note that strong net polarization consistently observed from SNe IIn, as a class, implies that CSM interaction is a dominanting influence on their polarization properties, and that an aspherical distribution of CSM is a commonplace feature of the stellar progenitors of these explosions.
14
3
1403.4240
We present spectropolarimetry of SN 2009ip throughout the evolution of its 2012 explosion. During the 2012a phase, when the spectrum exhibits broad P-Cygni lines, we measure a V-band polarization of P ≈ 0.9 per cent at a position angle of θ ≈ 166°, indicating substantial asphericity for the 2012a outflow. Near the subsequent peak of the 2012b phase, when the spectrum shows signs of intense interaction with circumstellar material (CSM), we measure P ≈ 1.7 per cent and θ ≈ 72°, indicating a separate component of polarization during 2012b, which exhibits a higher degree of asphericity than 2012a and an orthogonal axis of symmetry on the sky. Around 30 d past peak, coincident with a substantial bump in the declining light curve, we measure P ≈ 0.7 per cent and another significant shift in θ. At this point, broad photospheric lines have again become prominent and exhibit significant variations in P relative to the continuum, particularly He I/Na ID. By 60 d past peak, the continuum polarization has dropped below 0.2 per cent, probably declining towards a low value of interstellar polarization. The results are consistent with a scenario in which a prolate (possibly bipolar) explosion launched during the 2012a phase impacts an oblate (toroidal) distribution of CSM in 2012b. Previous calculations that assumed spherical symmetry for the CSM have substantially underestimated the required explosion energy, since only a small fraction of the SN ejecta appears to have participated in strong CSM interaction. A kinetic energy of ∼10<SUP>51</SUP> erg is difficult to avoid, supporting the interpretation that the 2012 outburst of SN 2009ip was the result of a core-collapse explosion.
false
[ "P ≈", "strong CSM interaction", "CSM", "interstellar polarization", "polarization", "substantial asphericity", "θ", "SN 2009ip", "intense interaction", "2012a", "2012b", "circumstellar material", "significant variations", "SN", "asphericity", "spherical symmetry", "the required explosion energy", "P", "broad photospheric lines", "symmetry" ]
14.089493
-0.529547
0
485723
[ "Skelton, Rosalind E.", "Whitaker, Katherine E.", "Momcheva, Ivelina G.", "Brammer, Gabriel B.", "van Dokkum, Pieter G.", "Labbé, Ivo", "Franx, Marijn", "van der Wel, Arjen", "Bezanson, Rachel", "Da Cunha, Elisabete", "Fumagalli, Mattia", "Förster Schreiber, Natascha", "Kriek, Mariska", "Leja, Joel", "Lundgren, Britt F.", "Magee, Daniel", "Marchesini, Danilo", "Maseda, Michael V.", "Nelson, Erica J.", "Oesch, Pascal", "Pacifici, Camilla", "Patel, Shannon G.", "Price, Sedona", "Rix, Hans-Walter", "Tal, Tomer", "Wake, David A.", "Wuyts, Stijn" ]
2014ApJS..214...24S
[ "3D-HST WFC3-selected Photometric Catalogs in the Five CANDELS/3D-HST Fields: Photometry, Photometric Redshifts, and Stellar Masses" ]
821
[ "South African Astronomical Observatory, PO Box 9, Observatory, Cape Town 7935, South Africa ; Department of Astronomy, Yale University, 260 Whitney Avenue, New Haven, CT 06511, USA;", "Astrophysics Science Division, Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Department of Astronomy, Yale University, 260 Whitney Avenue, New Haven, CT 06511, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Department of Astronomy, Yale University, 260 Whitney Avenue, New Haven, CT 06511, USA", "Leiden Observatory, Leiden University, Leiden, The Netherlands", "Leiden Observatory, Leiden University, Leiden, The Netherlands", "Max Planck Institute for Astronomy (MPIA), Königstuhl 17, D-69117, Heidelberg, Germany", "Department of Astronomy, Yale University, 260 Whitney Avenue, New Haven, CT 06511, USA; Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721, USA", "Max Planck Institute for Astronomy (MPIA), Königstuhl 17, D-69117, Heidelberg, Germany", "Leiden Observatory, Leiden University, Leiden, The Netherlands", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse, D-85748 Garching, Germany", "Astronomy Department, University of California, Berkeley, CA 94720, USA", "Department of Astronomy, Yale University, 260 Whitney Avenue, New Haven, CT 06511, USA", "Department of Astronomy, University of Wisconsin-Madison, 475 North Charter Street, Madison, WI 53706, USA", "Department of Astronomy &amp; Astrophysics, University of California, Santa Cruz, CA, USA", "Department of Physics and Astronomy, Tufts University, Medford, MA 02155, USA", "Max Planck Institute for Astronomy (MPIA), Königstuhl 17, D-69117, Heidelberg, Germany", "Department of Astronomy, Yale University, 260 Whitney Avenue, New Haven, CT 06511, USA", "Department of Astronomy, Yale University, 260 Whitney Avenue, New Haven, CT 06511, USA", "Yonsei University Observatory, Yonsei University, Seoul 120-749, Republic of Korea", "Carnegie Observatories, Pasadena, CA 91101, USA", "Astronomy Department, University of California, Berkeley, CA 94720, USA", "Max Planck Institute for Astronomy (MPIA), Königstuhl 17, D-69117, Heidelberg, Germany", "Department of Astronomy &amp; Astrophysics, University of California, Santa Cruz, CA, USA", "Department of Astronomy, University of Wisconsin-Madison, 475 North Charter Street, Madison, WI 53706, USA; Department of Physical Sciences, The Open University, Milton Keynes, MK7 6AA, UK", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse, D-85748 Garching, Germany" ]
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[ "astronomy" ]
68
[ "catalogs", "galaxies: evolution", "galaxies: general", "methods: data analysis", "techniques: photometric", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1088/0067-0049/214/2/24", "10.48550/arXiv.1403.3689" ]
1403
1403.3689.txt
\label{sec:intro} Large multi-wavelength photometric surveys have made it possible to study galaxy populations over most of cosmic history. Near-infrared selected samples have been used to trace the evolution of the stellar mass function \citep[e.g.,][]{Marchesini09, Perez-Gonzalez08}, the star formation--mass relation (e.g., \citealt{Whitaker12}), and the structural evolution of galaxies \citep[e.g.,][]{Franx08, Bell12, Wuyts12, vanderWel12}. Until recently most of these surveys relied on deep, wide-field imaging from ground-based telescopes (e.g., \citealt{Muzzin13a, Williams09}). The WFC3 camera on the \textit{Hubble Space Telescope} (\textit{HST}) has opened up the possibility to select and study galaxies at near-infrared wavelengths with excellent sensitivity and spatial resolution. Furthermore, the WFC3 grisms enable space-based near-IR slitless spectroscopy of all objects in the camera's field-of-view (see, e.g., \citealt{vanDokkumBrammer10}). The largest area WFC3 imaging survey done to date is the Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey (CANDELS, \citealt{Grogin11, Koekemoer11}), a 912-orbit Multi-Cycle Treasury imaging program (PIs: S.~Faber and H.~Ferguson). This survey encompasses five well-studied extragalactic fields: the All-wavelength Extended Groth Strip International Survey (AEGIS) field, the Cosmic Evolution Survey (COSMOS) field, the Great Observatories Origins Survey (GOODS) Northern and Southern fields (GOODS-North and GOODS-South) and the UKIRT InfraRed Deep Sky Surveys (UKIDSS) Ultra Deep Field (UDS). The coordinates of the five fields are given in Table~\ref{table:fields}. As these fields have been observed extensively over the past decade, the CANDELS imaging builds on a vast array of publicly available photometry at other wavelengths, ranging from the near-UV to the far-IR (see \citealt{Grogin11}). Building, in turn, on the CANDELS survey, we have undertaken a WFC3 spectroscopic survey in these same fields. 3D-HST is a 248-orbit \textit{HST} Treasury program (Programs 12177 and 12328; PI: P.\ van Dokkum) that uses the WFC3 G141 grism for slitless spectroscopy across $\sim700$ arcmin$^2$ of the sky, approximately 75\% of the CANDELS area (see Figure \ref{fig:fieldlayout}). This rich dataset is providing excellent redshifts and spatially-resolved spectral lines for thousands of galaxies in the key epoch $1 < z < 3$ (e.g., \citealt{Whitaker13, Nelson12, Brammer12b}). The survey is described in \citet{Brammer12a}. We targeted four of the five CANDELS fields: AEGIS, COSMOS, GOODS-South, and UDS. WFC3/G141 grism data in GOODS-North were already available from program GO-11600 (PI: B.\ Weiner); these data are incorporated in the 3D-HST analysis and data releases. The 3D-HST observations yield the following four types of data: WFC3 G141 grism observations; WFC3 F140W imaging for wavelength calibration of the spectra; parallel ACS G800L grism spectroscopy; and parallel F814W imaging. The scientific returns from CANDELS, 3D-HST and all other surveys in these five fields are maximized when the various datasets are combined in a homogeneous way, as it is often the combination of different kinds of data that provides new insight. To give just one example, \citet{Wuyts12} study the structure of galaxies (determined from \textit{HST} imaging) as a function of photometric redshift (determined from fits to multi-wavelength, broad-band spectral energy distributions, SEDs) and star formation rate (determined from SEDs and space-based infrared photometry). The interpretation of the data is also made easier when all information is used: it is much easier to correctly identify an emission line in a grism spectrum when the redshift range of the object is constrained by the available photometric information. The 3D-HST project has the aim of providing this homogenous combination of datasets in the five CANDELS/3D-HST fields. This undertaking has several linked aspects: \begin{enumerate} \item{We obtained and reduced the available HST/WFC3 imaging in the fields, using the same pixel scale and tangent point as those used by the CANDELS team. The WFC3 imaging includes the CANDELS data and also the Early Release Science data in GOODS-South and various other programs such as the HUDF09 Ultra Deep Field campaign.} \item{Source catalogs are created with SExtractor \citep{Bertin96}, detecting objects in deep combined F125W + F140W + F160W images.} \item{These source catalogs, along with the detection images, associated segmentation maps and PSFs, are used as the basis to measure photometric fluxes at wavelengths $0.3\,\mu$m -- $8\,\mu$m from a large array of publicly available imaging datasets. The resulting SEDs are of very high quality, particularly in fields with extensive optical and near-IR medium band photometry.} \item{Photometric redshifts, and redshift probability distributions, are estimated from the SEDs.} \item{Stellar population parameters are determined by fitting stellar population synthesis models to the SEDs, using the photometric redshifts as input.} \item{Mid- and far-IR photometry is obtained from Spitzer/MIPS and Herschel imaging. These data, combined with rest-frame UV emission measurements from the SEDs, are used to determine star formation rates of the galaxies.} \item{The set of images, PSFs, and catalogs is used to measure structural parameters of the objects in the WFC3 and ACS bands, following the methodology of \citet{vanderWel12}.} \item{The coordinates in the catalogs and segmentation maps are mapped back to the original (interlaced) coordinate system of the WFC3 and ACS grism data, and spectra are extracted for each object in the photometric catalog that is covered by the grism. No source matching is required, since each extracted spectrum is associated with a particular object in the photometric catalog. The photometric SED can be combined directly with the grism spectroscopy of each object for further analysis.} \item{The spectra and SEDs are fitted simultaneously, to measure redshifts and emission line fluxes.} \item{Parameters measured in steps 5, 6, and 7 are re-measured using the updated redshifts.} \end{enumerate} In this paper we describe steps 1--5 of the 3D-HST project; steps 6--10 will be described in future papers. As outlined above the photometric catalogs ultimately serve as input to the fits of the grism spectroscopy, but as we show here they constitute a formidable dataset in their own right. Furthermore, the majority of objects in the photometric catalogs are so faint that the grism does not provide useful additional information. We provide the homogenized set of imaging datasets that are used in this paper to the community, as well as the photometric catalogs and the EAZY and FAST fits to the photometry. The structure of this paper is as follows. In Section~\ref{sec:wfc3im} we describe the data reduction and mosaicking of the WFC3 detection images. Section~\ref{sec:otherdata} details the additional multi-wavelength data available for each field. Section~\ref{sec:phot} describes our photometric methods, accounting for differences in the depth and resolution of the data in different bands. We discuss the survey completeness in Section~\ref{sec:completeness}. We verify the quality and consistency of the catalogs in Section~\ref{sec:tests}. In Section~\ref{sec:eazy}, Section~\ref{sec:rfcols} and Section~\ref{sec:fast} we describe the photometric redshift, rest-frame color and stellar population parameter fits to the SEDs. Additional information on the PSFs and zero point offsets applied to the catalogs are provided in Appendix~\ref{app:psfs} and \ref{app:zps}. We present a comparison of our photometry with other available catalogs for each of the five fields in Appendix~\ref{app:comparisons}. We use the AB magnitude system throughout \citep{Oke71} and where necessary, a $\Lambda$CDM cosmology with $\Omega_M$=0.3, $\Omega_{\Lambda}$ = 0.7 and $H_0 = 70~$km s$^{-1}$ Mpc$^{-1}$. \begin{table*}[ht] \centering \caption{3D-HST Fields}\label{table:fields} \begin{tabular}{lllll} \hline \hline \noalign{\smallskip} & RA & Dec & Total area & Science Area \\ Field& (h m s) & (d m s) & (arcmin$^2$) & (arcmin$^2$) \\ \noalign{\smallskip} \hline \noalign{\smallskip} AEGIS &14 18 36.00 & +52 39 0.00 & 201 & 192.4 \\ COSMOS & 10 00 31.00 & +02 24 0.00 & 199 & 183.9 \\ GOODS-North & 12 35 54.98 & +62 11 51.3 & 164 & 157.8 \\ GOODS-South & 03 32 30.00 & -27 47 19.00 & 177 & 171.0 \\ UDS & 02 17 49.00 & -05 12 2.00 &201 & 191.2 \\ \noalign{\smallskip} \hline \noalign{\smallskip} \end{tabular} \end{table*}
In this paper we have presented the images and multi-wavelength photometric catalogs produced by the 3D-HST project for the five CANDELS/3D-HST extragalactic fields. The survey covers $\sim900$ arcmin$^2$ in the AEGIS, COSMOS, GOODS-North, GOODS-South and UDS fields with HST/WFC3 imaging and grism spectroscopy. The details of the WFC3 image reduction are given in \S\,~\ref{sec:wfc3reduction}. In addition to the new WFC3 data, we incorporated much of the available ground-based, Spitzer and HST/ACS data into the catalogs, using a total of 147 distinct data sets (see \S\,\ref{sec:otherdata} and Table~\ref{table:ancil_data}). We make all the images that have been used available on our website together with the catalogs. Each of the images is on the same astrometric system as the CANDELS WFC3 mosaics. We have applied consistent methodology to produce multi-wavelength catalogs for all five of the fields. The SExtractor software \citep{Bertin96} was used to detect sources on a noise-equalized combination of the F125W, F140W and F160W images. By using all three WFC3 bands, we exploited the maximum survey area and depth. As described in detail in \S\,\ref{sec:phot}, we measured the SEDs of objects using all the available ancillary data, carefully taking into account differences in image resolution (see \S\,\ref{sec:hstphot} and \ref{sec:lowresphot}). The results are consistent for all five of the fields and the total WFC3 magnitudes agree well with independently derived total magnitudes from morphological fitting (see \S\,\ref{sec:tests}). The resulting SEDs span from the $U$-band to 8$\mu$m and are of excellent quality, as demonstrated throughout the paper. We used the EAZY code \citep{Brammer08} to fit photometric redshifts and reach an NMAD scatter between the photometric and spectroscopic redshifts of $\lt 2.7\%$ with fewer than 5\% significant outliers in all fields. In the two fields where there is good medium band coverage (COSMOS and GOODS-S), the scatter is $\le 1\%$. We provide rest-frame colors based on the best-fitting EAZY templates, as well as stellar masses and stellar population parameters for all the galaxies based on fits to their SEDs. The CANDELS team has provided similar catalogs for two of the five fields discussed in this paper \citep{Guo13, Galametz13}, and we can expect future CANDELS releases of the other three fields. Our catalogs are complementary to these; we use slightly deeper detection images and a larger number of photometric filters, but these differences are probably not critical for most purposes. It will be very useful to have multiple "realizations" of the CANDELS datasets in the public domain, using independent reductions and methodology. In an Appendix we show band-by-band differences between the CANDELS catalogs and ours, and such comparisons provide much-needed estimates of systematic uncertainties in the various catalogs. As explained in the Introduction, in the context of the 3D-HST project, the work described in this paper merely concludes the first phase of an even more ambitious undertaking. The most innovative aspect of 3D-HST is the grism spectroscopy; future papers will describe these data and will quantify how they improve the measurement of redshifts, masses, and other parameters. \\ \\ \par We are grateful to the many colleagues who have provided public data and catalogs in the five fields described in this paper; high redshift galaxy science has thrived owing to this gracious mindset and the TACs and Observatory Directors who have encouraged this. We would like to thank the referee for a careful reading of the paper and useful suggestions. Formal acknowledgements follow below. RS acknowledges the support of the South African National Research Foundation through the Professional Development Programme Postdoctoral Fellowship. KW acknowledges support by an appointment to the NASA Postdoctoral Program at the Goddard Space Flight Center, administered by Oak Ridge Associated Universities through a contract with the National Aeronautics and Space Administration (NASA). DM acknowledges the support of the Research Corporation for Science Advancements Cottrell Scholarship. BL acknowledges support from the NSF Astronomy and Astrophsics Fellowship grant AST-1202963. We acknowledge support from ERC Advanced Grant HIGHZ\#227749, and a NWO Spinoza Grant. We thank the Lorentz Center for it support and hospitality. This work is based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the MAST Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. Observations associated with the following GO and GTO programs were used: 12063, 12440, 12442, 12443, 12444, 12445, 12060, 12061, 12062, 12064 (PI: Faber); 12177 and 12328 (PI: van Dokkum); 12461 and 12099 (PI: Riess); 11600 (PI: Weiner); 9425 and 9583 (PI: Giavalisco); 12190 (PI: Koekemoer); 11359 and 11360 (PI: OÕConnell); 11563 (PI: Illingworth). Based, in part, on data 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 NASA and was made possible by the generous financial support of the W. M. Keck Foundation. This work makes use of data obtained as part of the ESO/GOODS survey. Observations have been carried out using the Very Large Telescope at the ESO Paranal Observatory under Programme ID 168.A-0485. Based on observations made with ESO Telescopes at the La Silla or Paranal Observatories under programme number 168.A-0485. Based on data products from observations made with ESO Telescopes at the La Silla Paranal Observatory under ESO programme ID 179.A-2005 and on data products produced by TERAPIX and the Cambridge Astronomy Survey Unit on behalf of the UltraVISTA consortium. Based on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/IRFU, at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institut National des Science de l'Univers of the Centre National de la Recherche Scientifique (CNRS) of France, and the University of Hawaii. This work is based in part on data products produced at Terapix available at the Canadian Astronomy Data Centre as part of the Canada-France-Hawaii Telescope Legacy Survey, a collaborative project of NRC and CNRS. This work makes use of data products produced at TERAPIX, the WIRDS consortium, and the Canadian Astronomy Data Centre. We thank H. Hildebrandt for providing the CARS-reduced CFHTLS images. This research has made use of the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This study makes use of data from AEGIS, a multiwavelength sky survey conducted with the Chandra, GALEX, Hubble, Keck, CFHT, MMT, Subaru, Palomar, Spitzer, VLA, and other telescopes and supported in part by the NSF, NASA, and the STFC. This study makes use of data from the NEWFIRM Medium-Band Survey, a multi-wavelength survey conducted with the NEWFIRM instrument at the KPNO, supported in part by the NSF and NASA. This paper uses data products produced by the OIR Telescope Data Center, supported by the Smithsonian Astrophysical Observatory. This work is based on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under contract with NASA. Based in part on data from the UKIDSS Ultra Deep Survey (UDS) and the Spitzer Public Legacy Survey of the UKIDSS Ultra Deep Survey (SpUDS), a Cycle 4 Spitzer Legacy program. \vspace{0.2cm}\\ Author contributions: RS and KW led this part of the 3D-HST project and created the photometric catalogs described in this paper. RS, IM and PvD wrote the manuscript. IM is the 3D-HST project lead, reduced all the WFC3 imaging in the CANDELS fields, and was responsible for the integration of the data products. GB and IL developed many of the algorithms and techniques used in the creation of the catalogs. PvD is the PI of 3D-HST. The other authors provided help at various stages with the project and/or commented on the manuscript.
14
3
1403.3689
The 3D-HST and CANDELS programs have provided WFC3 and ACS spectroscopy and photometry over ≈900 arcmin<SUP>2</SUP> in five fields: AEGIS, COSMOS, GOODS-North, GOODS-South, and the UKIDSS UDS field. All these fields have a wealth of publicly available imaging data sets in addition to the Hubble Space Telescope (HST) data, which makes it possible to construct the spectral energy distributions (SEDs) of objects over a wide wavelength range. In this paper we describe a photometric analysis of the CANDELS and 3D-HST HST imaging and the ancillary imaging data at wavelengths 0.3-8 μm. Objects were selected in the WFC3 near-IR bands, and their SEDs were determined by carefully taking the effects of the point-spread function in each observation into account. A total of 147 distinct imaging data sets were used in the analysis. The photometry is made available in the form of six catalogs: one for each field, as well as a master catalog containing all objects in the entire survey. We also provide derived data products: photometric redshifts, determined with the EAZY code, and stellar population parameters determined with the FAST code. We make all the imaging data that were used in the analysis available, including our reductions of the WFC3 imaging in all five fields. 3D-HST is a spectroscopic survey with the WFC3 and ACS grisms, and the photometric catalogs presented here constitute a necessary first step in the analysis of these grism data. All the data presented in this paper are available through the 3D-HST Web site (http://3dhst.research.yale.edu).
false
[ "stellar population parameters", "WFC3", "account", "photometric redshifts", "publicly available imaging data sets", "the UKIDSS UDS field", "Objects", "objects", "derived data products", "wavelengths", "SEDs", "ACS", "the ancillary imaging data", "http://3dhst.research.yale.edu", "the FAST code", "≈900 arcmin", "147 distinct imaging data sets", "CANDELS programs", "the WFC3 imaging", "the EAZY code" ]
12.708032
6.111011
165
526921
[ "Magic, Zazralt", "Collet, Remo", "Asplund, Martin" ]
2014arXiv1403.6245M
[ "The Stagger-grid: A grid of 3D stellar atmosphere models - V. Fe line shapes, shifts and asymmetries" ]
10
[ "-", "-", "-" ]
[ "2015A&A...573A..90M", "2016AJ....152..167T", "2017A&A...597A...6N", "2017A&A...597A..52M", "2017A&A...605A..90D", "2017A&A...607A...6M", "2017A&A...607A.124M", "2018A&A...611A..11C", "2021A&A...649A..16D", "2021arXiv210406072M" ]
[ "astronomy" ]
7
[ "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1958ZA.....46..108B", "1965ApJ...142..841H", "1981A&A....96..345D", "1982A&A...107....1N", "1988ApJ...331..815M", "1989A&A...213..371S", "1990A&A...228..155N", "1991ApJ...370..295C", "1994A&A...291..635A", "1998ApJ...499..914S", "2000A&A...359..729A", "2000A&A...359..743A", "2000ApJ...536..465S", "2002ApJ...567..544A", "2003AN....324..399V", "2004ApJ...616..498C", "2005oasp.book.....G", "2008A&A...486..951G", "2008A&A...492..841R", "2009A&A...501.1087R", "2009ARA&A..47..481A", "2009ApJ...697.1032G", "2009ApJ...704L..66M", "2009LRSP....6....2N", "2010ApJ...725L.223R", "2011A&A...529A.158H", "2012JCoPh.231..919F", "2013A&A...557A..26M" ]
[ "10.48550/arXiv.1403.6245" ]
1403
1403.6245_arXiv.txt
} In recent years, capabilities for very high-resolution and very high signal-to-noise spectroscopical observations have raised the level of accuracy in stellar abundance analyses substantially \citep[e.g.,][]{Asplund:2005p7802,Melendez2009ApJ...704L..66M}. In addition truly large-scale, comprehensive high-resolution spectroscopic surveys are currently conducted, and further ones are planned. For the accurate interpretation of these sterling data, improved theoretical atmosphere models were also needed. Cool stars are characterized by convective envelopes that extend to the optical surface. The concomitant velocity field manifests itself in the stellar photosphere with a typical granulation pattern that imprints wavelength shifts and asymmetries in the observed spectral line profiles \citep[e.g.,][]{Dravins:1981p15588,Nordlund:2009p4109}. The strength of a spectral line depends mainly on the number of absorbers (atomic level population), therefore, it is very sensitive to the temperature due to exponential and power dependence of excitation and ionization equilibria (in LTE: $\propto e^{-\chi/kT}$; see \citealt{Gray2005oasp.book.....G} for more details on the theory of line formation in stellar atmospheres). In addition, owing to the presence of convective velocities, additional non-thermal broadening takes place in form of Doppler shifts. The wings of spectral lines are formed in deeper layers close to the continuum forming depth, while the cores are formed above in higher layers, therefore, the line profile samples different heights with distinctive physical conditions, in terms of velocity amplitudes, correlation between temperature and density inhomogeneities, asymmetries between regions with up- and downflowing material. This lead to characteristic asymmetries in the emergent intensity and flux line profiles \citep[e.g.,][]{Asplund:2000p20875}. Classical theoretical atmosphere models make use of several simplifications in order to facilitate calculations with the computing power at hand in the past \citep[e.g.,][]{Gustafsson:2008p3814,Cassisi:2004p1158}. The treatment of convection is a partially challenging part in modeling stellar atmospheres, since a complete theory of convection is still absent. In one-dimensional (1D) modeling, simplified treatments of convective energy transport have been adopted, such as the mixing-length theory \citep[MLT; see][]{BohmVitense:1958p4822,Henyey:1965p15592} or the full spectrum of turbulence model \citep{Canuto:1991p6553} with a priori unknown free-parameters that has to be calibrated by observations. Furthermore, for the 1D line formation calculations the lack of knowledge on the convective velocity fields is partially compensated by introducing two fudge parameters micro- and macroturbulence to account for convective broadening of spectral lines \citep[e.g.,][]{Gray2005oasp.book.....G}. For the precise modeling of realistic line profiles, including predicting their asymmetries, one has to rely on realistic three-dimensional (3D) atmosphere models, in which the convective velocity field emerges from first principles, i.e. from the solution of the hydrodynamic equations coupled with non-grey radiative transfer \citep[e.g.,][]{Nordlund:1982p6697,Steffen:1989p18861,Stein:1998p3801,Vogler:2003p11832,Nordlund:2009p4109,Freytag:2012p23073}. A major application for 3D radiative hydrodynamic (RHD) atmosphere models is the computation of synthetic full 3D line profiles or spectra as post-processing based on the former in order to derive accurate stellar parameters and abundances \citep{Asplund:2000p20875,Asplund:2000p20866,Asplund:2009p3308}. The 3D RHD models demonstrated their predictive capabilities powerfully in comparison with observed line profiles for several different stars. \citet{Asplund:2000p20875} found almost perfect agreement between observed solar iron line profiles and theoretical predictions without any trends in the derived abundances with line strength. Furthermore, comparisons of line shifts and asymmetries derived from high-resolution spectroscopical observations of different types of cool stars advocated additionally for the realistic nature of the theoretical 3D RHD models \citep[e.g.,][]{Nordlund:1990p6720,Atroshchenko:1994p14010,AllendePrieto:2002p22280,Ramirez2008A&A...492..841R,Ramirez:2009p10905,Ramirez2010ApJ...725L.223R,Gray2009ApJ...697.1032G}. With the present theoretical work we intend to tackle the following key question: how do line properties vary with stellar parameters? More specifically, we intend to analyze line shifts and asymmetries carefully for a selection of $\fei$ and $\feii$ lines, in order to better understand the variation of spectral line features with stellar parameters. Iron lines are often considered most useful for this purpose, since it is an abundant element with a very rich spectrum of energy levels and transitions with a variety of strengths and accurate atomic data at hand. In Sec. \ref{sec:Methods} we explain the methods for the computations of the 3D atmosphere models and line profiles. Subsequently, we discuss the properties of lines in terms of shape, strength, width, and depth (Sec. \ref{sec:Line-shape}, \ref{sub:Line-strength} and \ref{sub:Line-witdh-depth}), as well as the line asymmetries and shifts (Sec. \ref{sec:Line-asymmetry} and \ref{sec:Line-shift}). In Sec. \ref{sec:Height-of-line-formation}, we consider the physical conditions prevailing at the height of line forming region and link them with the findings on the line profiles. Finally in Sec. \ref{sec:Conclusions}, we summarize our results.
} We have explored the properties of synthetic spectral lines from neutral and singly ionized iron in late-type stars with the aid of 3D hydrodynamical model atmospheres. We have studied the variations with stellar parameters of aspects such as the strength, width, and depth of spectral lines, as well as line asymmetries and wavelength shifts. We have related such variations and the morphology of the asymmetries to the structural and thermal properties of the 3D models, with particular focus on velocity and temperature inhomogeneities and their correlation with depth in the stellar atmosphere. In Table \ref{tab:line_results} we list our results (line strength, shift, width, depth and bisectors). A possible application of the theoretical predictions of the line asymmetries can be the derivation of radial velocity and gravitational red-shift from high resolution observations by comparison with 3D line bisectors.
14
3
1403.6245
We present a theoretical study of the effects and signatures of realistic velocity field and atmospheric inhomogeneities associated with convective motions at the surface of cool late-type stars on the emergent profiles of iron spectral lines for a large range in stellar parameters. We compute 3D spectral line flux profiles under the assumption of local thermodynamic equilibrium (LTE) by employing state-of-the-art, time-dependent, 3D, radiative-hydrodynamical atmosphere models from the Stagger-grid. A set of 35 real unblended, optical FeI and FeII lines of varying excitation potential are considered. Additionally, fictitious Fe i and Fe ii lines (5000A and 0, 2, 4 eV) are used to construct general curves of growth and enable comparison of line profiles with the same line strength to illustrate systematical trends stemming from the intrinsic structural differences among 3D model atmospheres with different stellar parameters. Theoretical line shifts and bisectors are derived to analyze the shapes, shifts, and asymmetries imprinted in the full 3D line profiles emerging self-consistently from the convective simulations with velocity fields and atmospheric inhomogeneities. We find systematic variations in line strength, shift, width, and bisectors, that can be related to the respective physical conditions at the height of the line formation in the stellar atmospheric environment, in particular the amplitude of the vertical velocity field. Line shifts and asymmetries arise due to the presence of convective velocities and the granulation pattern that are ubiquitously found in observed stellar spectra of cool stars.
false
[ "line profiles", "3D spectral line flux profiles", "iron spectral lines", "Line shifts", "line strength", "convective velocities", "Theoretical line shifts", "velocity fields", "stellar parameters", "different stellar parameters", "realistic velocity field", "cool stars", "atmospheric inhomogeneities", "convective motions", "observed stellar spectra", "3D model", "the full 3D line profiles", "cool late-type stars", "local thermodynamic equilibrium", "the same line strength" ]
8.458656
10.832335
-1
621082
[ "Lattimer, James M.", "Steiner, Andrew W." ]
2014EPJA...50...40L
[ "Constraints on the symmetry energy using the mass-radius relation of neutron stars" ]
249
[ "Dept. of Physics &amp; Astronomy, Stony Brook University", "Institute for Nuclear Theory" ]
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"2022ResPh..4306037T", "2022SSPMA..52y2016X", "2022Symm...14..474L", "2022arXiv220302272B", "2022arXiv221010924C", "2022arXiv221116447B", "2023ApJ...950...77H", "2023JPhCS2586a2057B", "2023PhRvC.107d5803B", "2023PhRvC.107e5803Z", "2023PhRvC.108b4317Z", "2023PhRvC.108c5801S", "2023PhRvC.108e4305R", "2023PhRvC.108e5803R", "2023PhRvD.108a4018B", "2023PhRvD.108h3042L", "2023PhRvL.130k2701N", "2023Symm...15.1123K", "2023Univ....9..206Y", "2023arXiv230107884X", "2023arXiv230705018K", "2023arXiv231118819Y", "2024ApJ...966..184L", "2024EPJA...60...61S", "2024LRR....27....3K", "2024PhRvC.109b5805T", "2024PhRvC.109d5801L", "2024PhRvD.109h3007A", "2024PhRvD.109j3025I", "2024SSPMA..54v2011W", "2024arXiv240208835F", "2024arXiv240401989X", "2024arXiv240607396Z", "2024arXiv240610755L" ]
[ "astronomy", "physics" ]
6
[ "Nuclear Theory", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1966NucPh..81....1M", "1969AnPhy..55..395M", "1979PhRvC..19.1855R", "1989PhR...175..103L", "1990ARA&A..28..215D", "1990AnPhy.204..401M", "1994NuPhA.567..521K", "1994PhRvC..49.2118S", "1999PhRvL..82.3216K", "2000ApJ...542..914W", "2001ApJ...550..426L", "2001PhRvL..87h2501T", "2003NuPhA.727..233D", "2003NuPhA.729..337A", "2003PhRvC..67e4605C", "2004ApJ...613..512H", "2005PhR...411..325S", "2006Natur.441.1115O", "2007PhRvC..76b4606S", "2007PhRvC..76e1603K", "2008PhRvC..77b4317T", "2008PhRvC..77f1304T", "2009ApJ...693.1775O", "2009PhRvL.102l2701T", "2010ApJ...712..964G", "2010ApJ...719.1807G", "2010ApJ...722...33S", "2010Natur.467.1081D", "2010PhRvC..81e1303R", "2010PhRvC..82b4313K", "2010PhRvC..82b4321C", "2010PhRvC..82d4611Z", "2010PhRvD..82j1301O", "2010PhRvL.105p1102H", "2011ApJ...742..122S", "2011PhRvL.107f2502T", "2012ARNPS..62..485L", "2012ApJ...748....5O", "2012NuPhA.896...46F", "2012PhRvC..85c2801G", "2012PhRvL.108e2501M", "2012PhRvL.108h1102S", "2013ApJ...765L...1G", "2013ApJ...765L...5S", "2013ApJ...771...51L", "2013ApJ...772....7G", "2013ApJ...773...11H", "2013PhRvC..88b4316R", "2013Sci...340..448A", "2014ApJ...784..123L", "2014NuPhA.922....1D" ]
[ "10.1140/epja/i2014-14040-y", "10.48550/arXiv.1403.1186" ]
1403
1403.1186_arXiv.txt
\label{intro} Neutron stars are laboratories for the study of dense nuclear matter under conditions that are beyond those that can be achieved in experiments. The equation of state and internal compositions of large portions of neutron stars are poorly understood at present. However, there has been substantial recent progress in unraveling these mysteries. This progress has come from theoretical studies of nuclear and neutron matter, nuclear experiments, and astrophysical observations. The most profound aspect of the nuclear interaction for neutron stars, in many respects, concerns the nuclear symmetry energy which largely controls the composition and pressure of neutron-rich matter, and therefore, many aspects of neutron star structure such as the radius, moment of inertia and crustal properties. For most practical purposes, the interior of a neutron star can be divided into a dense core and a less-dense crust. The density of the core-crust boundary is believed to be near $n_s/2$, where $n_s\simeq0.16$ fm$^{-3}$ is the nuclear saturation density, with a weak dependence on the incompressibility and symmetry properties of bulk nuclear matter. While matter just below the neutron star crust is likely a uniform liquid of hadrons, electrons and muons, the crust itself is composed of an equilibrium mixture of dense nuclei and a neutron gas together with electrons. This division into two coexisting phases is a natural consequence of the fact that uniform nuclear matter at subnuclear densities, for large proton fractions, has negative pressure. For these densities, phase coexistence involves pressure and neutron and proton chemical potential equality in both phases, which together determine the relative concentrations of nuclei (the dense phase) and neutron gas (the less dense phase). It is important to point out that in the crust the pressure mostly originates from the degenerate relativistic electrons, for which the pressure is $p_e=\hbar cnx(3\pi^2nx)^{1/3}$, where $n$ is the baryon density and $x$ is the proton fraction (charge neutrality dictates that the number of electrons per baryon is also $x$). Baryon pressure originates from both nuclei and the neutron gas. However, the overall pressure of a nucleus must equal the neutron gas pressure, or it would expand or contract. This pressure remains very small until densities approach $n_s$. In fact, the dominant baryonic pressure results from the attractive Coulomb energy stemming from the Coulomb lattice, leading to a net negative pressure that, like $p_e$, scales as $n^{4/3}$. The ratio of the magnitudes of the lattice and electron pressures is only a few percent, however, so in spite of uncertainties regarding the nuclear force, the equation of state (at least, the pressure-density relation) in the crust is very well-understood. Matter in the interior of a neutron star, unlike that in laboratory nuclei, is very neutron rich. On timescales long compared to $\beta$ decay timescales of seconds, neutron star matter evolves into weak interaction equilibrium, or $\beta$-equilibrium, in which the total energy is at a minimum with respect to composition: \begin{equation}\label{beta} {\partial (E+E_e)\over\partial x}=\mu_p-\mu_n+\mu_e-(m_n-m_p)c^2=0, \end{equation} where $E$ is the baryon energy per baryon and $E_e$ is the electron energy per baryon. The $\mu$s are chemical potentials, which, for baryons, are measured with respect to their rest masses. Electrons are relativistic, and $\mu_e=\hbar c(3\pi^2nx)^{1/3}$. The energy of uniform hadronic matter, in its ground state, is essentially a function of baryon density ($n$), temperature ($T$), and composition, which is usually parameterized in terms of its charge fraction $x$. For baryonic matter composed solely of neutrons and protons, $x=n_p/(n_n+n_p)$. It is convenient to define the symmetry energy $S(n)$ as the difference between the energy per baryon of pure neutron matter ($x=0$) and symmetric nuclear matter ($x=1/2$). Since matter in neutron stars under nearly all conditions of interest here is highly degenerate, we only consider the case $T=0$. In most theoretical models of cold uniform nuclear matter, the energy at a given density can be well approximated by keeping only the first term of a quadratic expansion: \begin{equation}\label{energy} E(n, x) \simeq E(n, 1/2) + S_2(n)(1-2x)^2+\dots \end{equation} so that the symmetry energy $S(n)\simeq S_2(n)$. However, it has not been experimentally verified that quartic and higher-order terms are negligible. We will indicate where this neglect might have an appreciable effect. The symmetry energy is experimentally accessible from nuclear masses and other experiments such as dipole resonances and neutron skin thicknesses which sample matter near the nuclear saturation density $n_s$. It is therefore convenient to consider a Taylor expansion of $S_2$ near $n_s$: \begin{equation}\label{symmetry} S_2(n)\simeq S_v+{L\over3}(n-n_s)+{K_{\rm sym}\over18}(n-n_s)^2+\cdots \end{equation} which defines the symmetry parameters $S_v, L$ and $K_{\rm sym}$. From Equation (\ref{energy}), we now find that \begin{equation}\label{beta1} \mu_p-\mu_n={\partial E\over\partial x}=4S_2(n)(1-2x). \end{equation} The solution of Equation (\ref{beta}) at $n_s$ yields \begin{equation}\label{beta2} x\simeq{1\over3\pi^2n_s}\left({4S_v\over\hbar c}\right)^3\simeq0.04, \end{equation} i.e., neutron star matter is very nearly pure neutron matter. At higher densities, $x$ follows the behavior of $S_2(n)$. Below $n_s$, where nuclei exist, Equation (\ref{beta1}) shows that the neutron excess of the system, and individual nuclei, increases with density. The minimum value of $x$ in beta equilibrium generally occurs at the core-crust boundary just below $n_s$. For pure neutron matter at $n_s$ and in the quadratic approximation, the energy and pressure are given by \begin{eqnarray}\label{sat} E_N(n_s)&=&E(n,0)\simeq S_v+B,\cr p_N(n_s)&=&p(n_s,0)=n_s^2\left({\partial E\over\partial n} \right)_{n_s,x=0}\!\!\simeq{L\over3}n_s, \end{eqnarray} where $B=-E(n_s,1/2)\simeq16$ MeV is the binding energy of symmetric matter at the saturation density. For matter in $\beta$-equilibrium, it follows that \begin{equation}\label{sat1} p_\beta(n_s)\simeq{L\over3}n_s\left[1-\left({4S_v\over\hbar c} \right)^3{4-3S_v/L\over3\pi^2n_s}+\dots\right]. \end{equation} This important result shows that the pressure of matter at the nuclear saturation density can be expressed solely in terms of the standard symmetry parameters $S_v$ and $L$, in the quadratic approximation. The symmetry energy is not only important in determining the composition and pressure of matter in the interior, but also plays an important role in determining the overall structure of the star. Lattimer \& Prakash \cite{LP01} found that the neutron star radius $R$, for a given stellar mass $M$, is highly correlated with the neutron star matter pressure $p_\beta$ at densities in the vicinity of $n_s$. This relation can be expressed as \begin{equation}\label{radius} R_M=C(n,M)[p_\beta(n)/{\rm MeV~fm^{-3}}]^{1/4}, \end{equation} where $R_M$ is the radius of a star of mass $M$ and $C$ are coefficients that depend on the density and mass. The upper set in table~\ref{tab:1} shows the coefficients $C(n,1.4M_\odot)$ compiled from about 3 dozen equations of state for three densities, $n_s, 1.5n_s$ and $2n_s$. Lattimer \& Lim \cite{Lattimer13} re-analyzed this relation restricted to EOSs which could satisfy the constraint $\hat M=2.0M_\odot$ where $\hat M$ is the minimum value for the maximum neutron star mass given by the largest precisely measured neutron star mass. Currently, this determined from measurements of PSR J1614+2230 \cite{Demorest:2010}, with $M=1.97\pm0.04M_\odot$, and PSR J0348+0432 \cite{Antoniadis13}, with $M=2.01\pm0.04 M_\odot$. It is observed that the coefficients $C(n,M)$ become more accurate at higher densities, but since $p_\beta$ can be expressed relatively model-independently in terms of $S_v$ and $L$ at $n=n_s$, we can only usefully employ $C(n_s,1.4M_\odot)$ to relate neutron star radii to symmetry energy parameters. \begin{table} \begin{center} \caption{Coefficients $C(n,1.4M_\odot)$, in km, for the pressure-radius correlation. $\hat M$ is the minimum value for the maximum neutron star mass.\label{tab:1}} \begin{tabular}{l|ccc} \hline\noalign{\smallskip} $\hat M/M_\odot$&$n_s$ & $1.5n_s$ & $2n_s$\\ \noalign{\smallskip}\hline\noalign{\smallskip} 1.3&$9.30\pm0.58$&$6.99\pm0.30$&$5.72\pm0.25$\\ \noalign{\smallskip}\hline\noalign{\smallskip} 2.0&$9.52\pm0.49$ &$7.06\pm0.24$ &$5.68\pm0.14$\\ \noalign{\smallskip}\hline \end{tabular} \end{center} \end{table}
\label{sec:disc} A plethora of nuclear experimental data indicates that the symmetry energy parameters $S_v$ and $L$ are constrained to a greater degree than just a few years ago. Although these constraints have varying degrees of model dependence that need to be further explored, they are well-supported by studies of pure neutron matter, which can determine these parameters assuming that higher-than-quadratic terms in the symmetry energy expansion in neutron excess are ignored. It is expected that future theoretical studies of neutron matter with small proton concentrations will allow the validity of the quadratic expansion to be ascertained. From studies of solutions to the hydrostatic structure equations in general relativity \cite{LP01}, these symmetry energy restrictions and the quadratic approximation allow the radii of neutron stars to be determined to about 10\% accuracy \cite{Lattimer13}. For the experimental constraints studied here, the deduced radius of $1.4M_\odot$ neutron stars is $R_{1.4}=12.1\pm1.1$ km. Neutron matter studies suggest slightly smaller values by about 0.2 km. In comparison, the astrophysical determination of individual neutron star radii have much less precision. Nevertheless, Bayesian studies (cf., \cite{SLB10,SLB13}) of the ensemble of individual sources for which both mass and radius information is available, imply typical radii (i.e., for $1.2-1.8 M_\odot$ stars) in the range 11.2 -- 12.8 km. There is emerging an important interplay between the nuclear physics and the astronomical observations: {\em we find a concordance between the observations and the nuclear experiments.} With almost any reasonable assumptions regarding the nature of the EOS at high densities and the parameters of models for shorter PRE X-ray bursts and QLMXBs, the powerful constraints of causality, observations of 2$M_{\odot}$ neutron stars, and the existence of a nuclear neutron star crust, lead to $M-R$ curves which are nearly vertical and radii for moderate-mass neutron stars that are compatible with nuclear data and theoretical studies of neutron matter. Thus, neutron star mass and radius observations are clearly beginning to make quantitative constraints on both the EOS and the parameter $L$ which describes the density dependence of the symmetry energy. Two major classes of neutron star observations have provided important constraints: PRE X-ray bursts and the surface emission of QLMXBs. In both of these classes of neutron star observations, the theoretical models which interpret X-ray photons and produce the inferred neutron star mass and radius are an important source of uncertainty. PRE X-ray bursts are interpreted as resulting from the vertical motion of the photosphere. Assumptions about the position of the photosphere at touchdown can change radius estimates by about 2 km. If the photosphere of PRE X-ray burst neutron stars is redshifted at touchdown, we find that the observed fluxes and normalizations tend to be inconsistent with the model, judging from the small number of Monte Carlo trials over the observed uncertainty ranges of touchdown fluxes, distances, and normalizations that result. In addition, the 95\% confidence radius range from Ref.~\cite{Ozel12,Guver13,Ozel:2010}, which comes from PRE sources alone and assumes the photosphere at touchdown is at the stellar surface, i.e., $z_{\mathrm{ph}}=z$, is also incompatible with nuclear experiment, as seen in Figure \ref{fid}. However, these are not the only difficulties surrounding the interpretation of PRE X-ray bursts, and color correction factors and composition are also important uncertain parameters. For example, Suleimanov et al. have argued \cite{Suleimanov:2011} that the short PRE bursts studied by Ozel et al. and in this contribution might have significant disk absorption and $f_c$ evolution during the burst that would dramatically increase the inferred radii. Ref. \cite{Suleimanov:2011} instead studied longer PRE bursts and found radii in excess of 13.9 km to 90\% confidence (Figure \ref{fid}), assuming stellar masses less than $2.3M_\odot$. Importantly, both the ranges suggested by Ozel et al. and Suleimanov et al. are inconsistent with nuclear systematics. In the case of QLMXBs, there is no photospheric dynamics to complicate the interpretation of the neutron star atmosphere, but the composition of the atmosphere and the magnitude of X-ray absorption between us and the source are both major uncertainties. Differences of assumed X-ray absorption magnitudes result in both larger and smaller radii. If the hydrogen column densities are assumed to be those obtained from self-consistent fitting of X-ray spectra \cite{Guillot13}, in some cases the observed neutron stars are (i) too small to satisfy causality limits, and (ii) too large to be consistent with the available nuclear data and any reasonable neutron star model. On the basis of our Bayesian model, however, we conclude that, on average, $N_H$ values from Ref.~\cite{Dickey90} are statistically favored in comparison to those obtained from self-consistently fitting \cite{Guillot13} the X-ray spectra. The alternative $N_H$ values also lead to a more uniform distribution of masses and radii among the sources. The radius range deduced by G13 in their joint study in which it is assumed that all neutron stars have the same radius is inconsistent with both our results from the joint study of PRE bursts and QLMXBs and with inferences from nuclear experiments to 90\% confidence. The large degree of model-dependence in interpreting astronomical observations suggests more sophisticated modeling is in order. It will be necessary to model PRE bursts using hydrodynamical radiation transport simulations to fit the overall light-curve behavior to fully resolve the discrepancies and to provide reliable $M$ and $R$ estimates. Similarly, for the QLMXBs, there is a clear necessity of obtaining further observations for fixing the interstellar X-ray absorption for QLMXBs. Moreover, there is evidence that models of QLMXBs allowing for the possibility of He as well as H atmospheres are favored, a question which further observations may also be able to decide.
14
3
1403.1186
The nuclear symmetry energy is intimately connected with nuclear astrophysics. This contribution focuses on the estimation of the symmetry energy from experiment and how it is related to the structure of neutron stars. The most important connection is between the radii of neutron stars and the pressure of neutron star matter in the vicinity of the nuclear saturation density n<SUB>s</SUB>. This pressure is essentially controlled by the nuclear symmetry energy parameters S<SUB>v</SUB> and L , the first two coefficients of a Taylor expansion of the symmetry energy around n<SUB>s</SUB>. We discuss constraints on these parameters that can be found from nuclear experiments. We demonstrate that these constraints are largely model-independent by deriving them qualitatively from a simple nuclear model. We also summarize how recent theoretical studies of pure neutron matter can reinforce these constraints. To date, several different astrophysical measurements of neutron star radii have been attempted. Attention is focused on photospheric radius expansion bursts and on thermal emissions from quiescent low-mass X-ray binaries. While none of these observations can, at the present time, determine individual neutron star radii to better than 20% accuracy, the body of observations can be used with Bayesian techniques to effectively constrain them to higher precision. These techniques invert the structure equations and obtain estimates of the pressure-density relation of neutron star matter, not only near n<SUB>s</SUB>, but up to the highest densities found in neutron star interiors. The estimates we derive for neutron star radii are in concordance with predictions from nuclear experiment and theory.
false
[ "neutron star matter", "neutron star radii", "individual neutron star radii", "nuclear experiment", "nuclear experiments", "neutron star interiors", "neutron stars", "neutron star", "nuclear astrophysics", "pure neutron matter", "higher precision", "SUB", "s</SUB", "photospheric radius expansion bursts", "experiment", "the nuclear saturation density", "The nuclear symmetry energy", "quiescent low-mass X-ray binaries", "Bayesian techniques", "constraints" ]
4.401753
1.964397
71
437446
[ "Agúndez, M.", "Biver, N.", "Santos-Sanz, P.", "Bockelée-Morvan, D.", "Moreno, R." ]
2014A&A...564L...2A
[ "Molecular observations of comets C/2012 S1 (ISON) and C/2013 R1 (Lovejoy): HNC/HCN ratios and upper limits to PH<SUB>3</SUB>" ]
25
[ "Univ. Bordeaux, LAB, UMR 5804, 33270, Floirac, France", "LESIA, Observatoire de Paris, CNRS, UPMC, Université Paris-Diderot, 5 place Jules Janssen, 92195, Meudon, France", "Instituto de Astrofísica de Andalucía - CSIC, Glorieta de la Astronomía s/n, 18008, Granada, Spain", "LESIA, Observatoire de Paris, CNRS, UPMC, Université Paris-Diderot, 5 place Jules Janssen, 92195, Meudon, France", "LESIA, Observatoire de Paris, CNRS, UPMC, Université Paris-Diderot, 5 place Jules Janssen, 92195, Meudon, France" ]
[ "2014A&A...566L...5B", "2014A&A...567L...1C", "2014ApJ...790L..27A", "2014ApJ...791..118M", "2014ApJ...792L...2C", "2014ApJ...796L...6B", "2015A&A...575A..52S", "2015A&A...584A.121O", "2015AJ....149...19K", "2016A&A...588A..72W", "2016ApJ...820...34D", "2016IAUFM..29A.228B", "2016Icar..266..152D", "2017A&A...604A.131B", "2017ApJ...838...33A", "2017ApJ...838..147C", "2017MNRAS.464.4282A", "2019ApJ...870L..26C", "2020A&A...635A..31M", "2020ApJ...894..132M", "2021ApJ...907...51N", "2021FrASS...8...43Z", "2021JQSRT.27207795U", "2023JPCA..127.1000B", "2023arXiv230108492C" ]
[ "astronomy" ]
9
[ "comets: general", "comets: individual: C/2012 S1 (ISON)", "comets:", "individual: C/2013 R1 (Lovejoy)", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1994Icar..107..164B", "1996Natur.383..418I", "1997Sci...275.1915B", "1997ssis.book.....K", "1998ApJ...497L.117L", "1998Natur.393..547I", "2000A&A...353.1101B", "2001MNRAS.323...84R", "2002EM&P...89...53B", "2004A&A...418.1141C", "2004ApJ...615.1054M", "2005JMoSp.229..170I", "2006A&A...449.1255B", "2008ApJ...675..931L", "2008M&PS...43..399G", "2009A&A...495..975A", "2009ARA&A..47..481A", "2009EM&P..105..267C", "2011IAUS..280..261B", "2012A&A...538A..89C", "2012A&A...539A..68B", "2012A&A...542L...3K", "2012CBET.3238....1N", "2013A&A...560A.101O", "2013ApJ...779L...3L", "2013CBET.3433....1O", "2013CBET.3649....1G", "2013CBET.3711....1C", "2013CBET.3715....1B", "2013IAUC.9266....1C", "2013JQSRT.129..158V", "2014ApJ...782L..37K" ]
[ "10.1051/0004-6361/201423639", "10.48550/arXiv.1403.0463" ]
1403
1403.0463_arXiv.txt
Radio spectroscopic observations of comets during their visit to the inner solar system have allowed to detect a wide variety of molecules in their coma (e.g., \cite{boc2002} 2002). These observations have provided significant constraints on the chemical nature of comets coming from the two main solar system reservoirs, the Oort cloud and the Kuiper belt, whose composition is expected to reflect to some extent that of the regions of the protosolar nebula where they were once formed. Two bright comets coming from the Oort cloud approached the Sun in late 2013, allowing us to perform sensitive radio spectroscopic observations and to probe their volatile content. C/2012 S1 (ISON) --hereafter ISON-- was discovered on September 2012 at 6.3 au from the Sun using a 0.4-m telescope of the International Scientific Optical Network (\cite{nev2012} 2012). It is a sungrazing comet, which at perihelion, on 2013 November 28.8 UT, passed at just 0.012 au from the Sun (MPEC 2013-Q27). Its orbital elements are consistent with a dynamically new comet, with fresh ices not previously irradiated by sunlight. A worldwide observational campaign has extensively followed this comet from heliocentric distances beyond 4 au (\cite{oro2013} 2013; \cite{li2013} 2013) to disappearance around perihelion (\cite{kni2014} 2014). C/2013 R1 (Lovejoy) --hereafter Lovejoy-- was discovered in September 2013 at $r_h$ = 1.94 au by Terry Lovejoy using a 0.2-m telescope (\cite{gui2013} 2013). This comet reached perihelion on 2013 December 22.7 UT. According to its orbital elements (MPEC 2014-D13), this is not its first perihelion passage. In this Letter we report IRAM 30m spectroscopic observations of the comets ISON and Lovejoy carried out when they were at heliocentric distances of $\sim$0.6 and $\sim$1 au, respectively.
We have carried out IRAM 30m observations of the comets C/2012 S1 (ISON) and C/2013 R1 (Lovejoy) at heliocentric distances of $\sim$0.6 and $\sim$1 au, respectively. We detected HCN, HNC, and CH$_3$OH in both comets, plus the ion HCO$^+$ in ISON and a few weak unidentified lines in Lovejoy, one of which might be assigned to CH$_3$NH$_2$. A tenfold enhancement of the HCN $J$ = 3-2 line was observed in comet ISON within less than 24 h on November 14, indicating an outburst of activity whose origin could be related to nucleus splitting. The large number of CH$_3$OH lines observed was used to derive kinetic temperatures in the coma of 90 and 60 K in ISON and Lovejoy, respectively. The HNC/HCN ratios derived, 0.18 in ISON and 0.05 in Lovejoy, are similar to those found in most previous comets and are consistent with an enhancement of HNC as the comet approaches the Sun. PH$_3$ was also searched for unsuccessfully in both comets so that only upper limits to the PH$_3$/H$_2$O ratio 4-10 times above the solar P/O elemental ratio were derived.
14
3
1403.0463
We present molecular observations carried out with the IRAM 30 m telescope at wavelengths around 1.15 mm towards the Oort cloud comets C/2012 S1 (ISON) and C/2013 R1 (Lovejoy) when they were at ~0.6 and ~1 au, respectively, from the Sun. We detect HCN, HNC, and CH<SUB>3</SUB>OH in both comets, together with the ion HCO<SUP>+</SUP> in comet ISON and a few weak unidentified lines in comet Lovejoy, one of which might be assigned to methylamine (CH<SUB>3</SUB>NH<SUB>2</SUB>). The monitoring of the HCN J = 3-2 line showed a tenfold enhancement in comet ISON on November 14.4 UT due to an outburst of activity whose exact origin is unknown, although it might be related to some break-up of the nucleus. The set of CH<SUB>3</SUB>OH lines observed was used to derive the kinetic temperature in the coma, 90 K in comet ISON and 60 K in comet Lovejoy. The HNC/HCN ratios derived, 0.18 in ISON and 0.05 in Lovejoy, are similar to those found in most previous comets and are consistent with an enhancement of HNC as the comet approaches the Sun. Phosphine (PH<SUB>3</SUB>) was also searched for unsuccessfully in both comets through its fundamental 1<SUB>0</SUB>-0<SUB>0</SUB> transition, and 3σ upper limits corresponding to PH<SUB>3</SUB>/H<SUB>2</SUB>O ratios 4-10 times above the solar P/O elemental ratio were derived. <P />Based on observations carried out with the IRAM 30 m telescope. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).Tables 3 and 4 are available online at <A href="http://www.aanda.org/10.1051/0004-6361/201423639/olm">http://www.aanda.org</A> and also 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/L2">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/564/L2</A>
false
[ "comet ISON", "comet Lovejoy", "ISON", "Lovejoy", "anonymous ftp", "Sun", "the Oort cloud comets", "most previous comets", "upper limits", "CH<SUB>3</SUB>OH lines", "both comets", "the comet", "a few weak unidentified lines", "PH", "HNC", "activity", "P", "November", "CDS", "Oort" ]
8.938463
16.961338
92
1922567
[ "Prša, Andrej", "Robin, Annie", "Barclay, Thomas" ]
2015IJAsB..14..165P
[ "Stellar statistics along the ecliptic and the impact on the K2 mission concept" ]
3
[ "Department of Astrophysics and Planetary Science, Villanova University, 800 E Lancaster Ave, Villanova, PA 19085, USA", "Institute Utinam, CNRS UMR6213, Université de Franche-Comté, OSU THETA de Franche-Comté-Bourgogne, Besançon, France", "NASA Ames Research Center, M/S 244-30, Moffett Field, CA 94035, USA" ]
[ "2015RAA....15.1945S", "2016AJ....151...68K", "2018exha.book.....P" ]
[ "astronomy" ]
4
[ "extrasolar planets", "Kepler", "stellar populations", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1989ApJ...345..245C", "1997A&A...324...65N", "1999ASPC..192.....H", "1999BaltA...8..147V", "2003A&A...409..523R", "2003AdSpR..31..345B", "2005A&A...436..895G", "2006A&A...453..635M", "2009A&A...495..819R", "2009A&A...503L..21M", "2009ExA....23..329C", "2010ApJ...725.2166H", "2010ApJS..190....1R", "2010Sci...327..977B", "2010Sci...330..653H", "2011AJ....141...83P", "2011AJ....142..160S", "2011ApJ...727...24T", "2011ApJ...730....3S", "2011ApJ...738..170M", "2011Sci...332..213C", "2012A&A...538A.106R", "2012Ap&SS.341...31D", "2012ApJS..201...15H", "2013ApJ...766...81F", "2013ApJ...767...82G", "2013ApJ...767..127H", "2013ApJ...778...53D", "2013ApJS..204...24B", "2013PhDT.........6P", "2014A&A...564A.102C", "2014A&A...569A..13R", "2014ApJ...784...45R", "2014ApJS..210...19B", "2014PASP..126..398H" ]
[ "10.1017/S1473550414000329", "10.48550/arXiv.1403.5888" ]
1403
1403.5888_arXiv.txt
Space-borne missions CoRoT \citep{baglin2003} and \kepler~\citep{borucki2010} revolutionized stellar and planetary physics by providing us with ultra-precise photometric data of $\sim$30\,ppm. The duty cycle of observations exceeded 90\% and, for the first time, we had a nearly uninterrupted time coverage of over 200,000 objects. The two missions predominantly catered for two overlapping communities, the exoplanetary science and asteroseismology. The detection of extra-solar planets using the transit method boosted their numbers from dozens to nearly 5000 \citep{batalha2013, burke2014, rowe2014}, and the number is growing still as data are being mined. At the same time, asteroseismology witnessed an explosion in novel techniques and exciting new results \citep{chaplin2011}, ranging from main sequence B stars \citep{papics2013} to solar-like oscillations in red giants \citep{gaulme2013}. The overlap between the fields is provided by using asteroseismic techniques that provide fundamental properties of planet candidate host stars, which in turn enables exoplanet researchers to obtain precise fundamental properties of planets \citep{huber2013}. Unfortunately, CoRoT suffered from a computer failure in November 2012 and attempts to restore it ceased in June 2013. \kepler~lost a second reaction wheel in May 2013, causing the telescope to go to a prolonged safe-mode; attempts to bring Kepler back to operational state ceased in July 2013. However, this did not imply that \kepler~is retired; a proposal to use solar photon pressure to balance a 2-wheel \kepler~satellite enabled a continued operation. For this balancing to work, the telescope must point approximately in the direction of the ecliptic, so the observations of the initial \kepler~field are no longer possible. The spacecraft can hold pointing within $+50^\circ$ and $-30^\circ$ of its velocity vector in the orbital plane (approximately the ecliptic), with the two remaining reaction wheels holding the cross-boresight pointing steady. The spacecraft roll is minimized through regular thruster firing windows. The satellite can remain stable in roll for up to 85 days with a fuel budget that allows for a 2-3 year mission duration. A new mission concept, K2 \citep{howell2014}, builds on this engineering constraint. The science case arose from the community response to the whitepaper call for the repurposed mission and the concept was submitted to NASA HQ for the 2014 Senior Review (pending at the time of this writing). The first engineering observations utilizing the K2 mission design concept were obtained in October 2013 and the first full campaign-length test began in March 2014. This field lies in the direction near the galactic anti-center ($\ra = \hms{6}{33}{11.1}$, $\dec = \dms{21}{35}{16}$) and includes M35 and NGC 2158. The subsequent fields will be observed for 83 days; the duration is limited by solar illumination. If selected, K2 will observe upwards of 40,000-80,000 targets over the first year in 4 distinct fields. \kepler~is a 0.95-m telescope with a 105-deg${}^2$ field of view. The early science commissioning run from October 2013 to February 2014 showed that the precision of K2 photometry for a $V=12$ star is $\sim$400\,ppm for the 30-min long cadence exposure and $\sim$80\,ppm for an integrated 6-hr exposure. The precision primarily depends on the spacecraft attitude jitter; the point spread function (PSF) of the K2 field is within 5\% of the original \kepler~field, and the degradation in precision due to a solar-induced drift is $\sim$4-fold \citep{howell2014}. Early science demonstration for WASP-28, a hot jupiter orbiting a Sun-like star in a 3.4-day orbit, corresponds to a 6-hr integrated noise level of $84$\,ppm. This work employs the updated Besan\c con model of the Galaxy \citep[Robin et al.~2014, submitted]{robin2003} to simulate stellar populations along the ecliptic. The K2 mission has the potential observe $\sim$250,000 stars, and selecting targets hinges crucially on the representative population within each K2 field. The goal of this paper is to study the bulk properties and to serve as a guide to stellar populations along the ecliptic. This information can be used to better understand different populations from which the K2 targets are drawn, and to enable debiasing of any results that stem from statistical analyses of K2 campaigns.
The K2 mission concept promises to yield invaluable data similar in nature to the original \kepler~mission, and akin to the upcoming Transiting Exoplanet Survey Satellite (TESS; \citealt{ricker2010}). With $\sim$85 days on a single field, K2 will probe inherently different stellar populations and, contingent on NASA HQ approval and continued funding for 2.5 years, provide photometric coverage of over 250 thousand stars. This paper provides a guide to expected stellar populations, crowding and planetary occurrence rates along the ecliptic based on the improved Besan\c con model simulations. Understanding stellar binarity and multiplicity is the next step in the study of K2 campaigns. From \kepler~observations of 2615 eclipsing and ellipsoidal binary stars \citep{prsa2011, slawson2011} that are essentially complete to $P \sim 500$ days, we can derive the underlying orbital period and eccentricity distributions. We do this by correcting for the bias using Bayesian methods outlined in \citet{hogg2010}. From the underlying distributions we simulate binary and multiple stars by applying the observed occurrence rates from \citet{raghavan2010} to the Besan\c con sample of stars that are grouped into multiple systems under the constraints of coevality and equal metallicity. These systems are then statistically examined for eclipses and eclipse timing variations. This in-depth analysis requires a subtantial discussion that is beyond the scope of the present paper. Inherently different stellar (and planetary) populations along the ecliptic provide us with an opportunity to study population differences as a function of galactic latitude. Table \ref{tab:bulk} lists the expected numbers of main sequence stars and giants for each campaign, attesting to the variety of objects for which K2 will provide ultraprecise, long temporal baseline photometry. In combination with Gaia \citep{debruijne2012} that has recently seen first light, TESS that is scheduled to launch in 2017, and Plato \citep{catala2009} that has been selected as the third ETS medium-class science mission, the K2 dataset will be a gold mine for stellar and planetary astrophysics.
14
3
1403.5888
K2 is the mission concept for a repurposed Kepler mission that uses two reaction wheels to maintain the satellite attitude and provide ~81 days of coverage for ten 105 deg<SUP>2</SUP> fields along the ecliptic in the first 2.5 years of operation. We examine stellar populations based on the updated Besançon model of the Galaxy, comment on the general properties for the entire ecliptic plane, and provide stellar occurrence rates in the first six tentative K2 campaigns grouped by spectral type and luminosity class. For each campaign we distinguish between main the sequence stars and giants, and provide their density profile as a function of galactic latitude. We introduce the crowding metric that serves for optimized target selection across the campaigns. For all main sequence stars we compute the expected planetary occurrence rates for three planet sizes: 2-4, 4-8 and 8-32 R <SUB>⊕</SUB> with orbital periods up to 50 days. In conjunction with Gaia and the upcoming Transiting Exoplanet Survey Satellite and Plato missions, K2 will become a gold mine for stellar and planetary astrophysics.
false
[ "stellar occurrence rates", "~81 days", "> fields", "operation", "galactic latitude", "first", "stellar populations", "orbital periods", "SUP>2</SUP", "deg", "coverage", "Transiting Exoplanet Survey Satellite", "optimized target selection", "spectral type and luminosity class", "the expected planetary occurrence rates", "K2", "stellar and planetary astrophysics", "the first six tentative K2 campaigns", "Galaxy", "the entire ecliptic plane" ]
6.947182
13.17522
-1
256584
[ "Khoperskov, S.", "Shchekinov, Y." ]
2013lcdu.confE..60K
[ "Transport of the charged dust grains to the galactic halo" ]
3
[ "-", "-" ]
[ "2014Sci...345..791K", "2015MmSAI..86..541Z", "2017PhyU...60..961S" ]
[ "astronomy" ]
2
[ "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
null
[ "10.22323/1.207.0060", "10.48550/arXiv.1403.7075" ]
1403
1403.7075_arXiv.txt
\vskip -0.021\vsize The dominant sources of dust in the Milky Way are the AGB stars~\cite{2009ApJ...698.1136G}, supernovae and young stellar objects and cool dense ISM regions~\cite{1998ApJ...501..643D}. Stellar winds and jets from YSO supply dust into the nearest interstellar medium where it is supposed to mix with ambient gas. SNe ejecta are also thought to be an efficient source of dust, though it is unclear what fraction of dust survives in the reverse shocks of SNe remnants~\cite{2012ApJ...748...12S}. However, it is absolutely clear that the sources of dust concentrate basically within the galactic stellar disc. At the same time, there are copious evidences of dust extending in the vertical direction up to tens of the scale height of the stellar thin disc~\cite{2006A&A...445..123I,2007A&A...471L...1K,1997AJ....114.2463H,2009AJ....137.3000H}. In the radial direction dust may occupy at least twice as large a disc as the stellar one. In addition, more recently evidences of dust present in the intergalactic medium have also appeared in the literature~\cite{2010MNRAS.405.1025M}. \begin{figure}[!b] \begin{center} \includegraphics[width=0.49\hsize]{khoperskov_fig01a.eps} \includegraphics[width=0.49\hsize]{khoperskov_fig01b.eps} \end{center} \caption{Contributions to the vertical acceleration of dust grains of radius $a=10^{-4}$~cm (left panel) and $a=10^{-6}$~cm~(right panel): gravity from the disc (dark blue), from the disc and the halo~(light blue); radiation pressure for $1$ (green) and $10$~Habing fluxes (orange); drag forces for a grain moving through the gaseous halo with a velocity $v = 10$~\kmps~(red) and $100$~\kmps~(brown).}\label{fig::accel} \end{figure} In this contribution we briefly describe our results of simulations of dust transport within a 3D N-body/hydrodynamical framework on scales of the galactic halo. Aiming to understand of how the galactic halo and galaxy outskirts are enriched with dust we developed a numerical scenario of dust driven by a combined action of stellar radiation pressure, multiple supernova explosions and Lorenz forces in a magnetized gravitationally stratified ISM.
\vskip -0.021\vsize Multiple supernovae explosions in the galactic thin stellar disc drive gas and dust into the halo. A large amount of dust can be transported upto $10$~kpc in $10-20$~Myr. Ram pressure from supernovae explosions reorganizes magnetic field structure such that it efficiently accelerates dust grains. The dust outflow rate strongly depends on the plasma-$\beta$ parameter and gas density in the disc. The estimated dust mass loss rate is in the range from $10^{-3}$~\Msun~yr$^{-1}$ up to $3\times 10^{-2}$~\Msun~yr$^{-1}$. Such a high rate is kept though on short time scales $2-10$~Myr.
14
3
1403.7075
We develop a 3D dynamical model of dust outflows from galactic discs. The outflows are initiated by multiple SN explosions in a magnetized interstellar medium (ISM) with a gravitationally stratified density distribution. Dust grains are treated as particles in cells interacting collisionally with gas, and forced by stellar radiation of the disc and Lorenz force. We show that magnetic field plays a crucial role in accelerating the charged dust grains and expelling them out of the disc: in 10-20 Myr they can be elevated at distances up to 10 kpc above the galactic plane. The dust-to-gas ratio in the outflowing medium varies in the range 5 · 10-4 - 5 · 10-2 along the vertical stream. Overall the dust mass loss rate depends on the parameters of ISM and may reach up to 3 × 10-2 M⊙ yr-1 .
false
[ "galactic discs", "Lorenz force", "Dust grains", "dust outflows", "stellar radiation", "ISM", "multiple SN explosions", "distances", "Lorenz", "the galactic plane", "gas", "the charged dust grains", "the outflowing medium varies", "the dust mass loss rate", "cells", "magnetic field", "the vertical stream", "a magnetized interstellar medium (ISM", "SN", "particles" ]
12.153652
8.29449
-1
621860
[ "Huang, Qing-Guo" ]
2014EPJC...74.2964H
[ "An analytic calculation of the growth index for dark energy model" ]
11
[ "State Key Laboratory of Theoretical Physics, Institute of Theoretical Physics, Chinese Academy of Science" ]
[ "2014PhRvD..90d3513R", "2019EPJP..134..318L", "2019IJMPD..2850132L", "2019PhLB..795..129L", "2019PhRvD.100h3503C", "2020PhLB..81135985L", "2021EPJP..136..883B", "2022ApJ...934...13K", "2022IJMPD..3150117L", "2022arXiv220306741S", "2023FoPh...53....5L" ]
[ "astronomy", "physics" ]
9
[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Theory" ]
[ "1980PhLB...91...99S", "1998ApJ...508..483W", "2000PhLB..485..208D", "2002PhRvD..65d4023D", "2005PhRvD..72d3529L", "2006PhRvD..73l3504Z", "2007APh....28..481L", "2007JETPL..86..157S", "2007PhLB..654....7A", "2007PhRvD..75h3504A", "2007PhRvD..76b3514T", "2007PhRvD..76f4004H", "2008PhLB..660..439P", "2008PhLB..664....1W", "2008PhRvD..77b3507T", "2008PhRvD..77d6009C", "2008PhRvD..77l3515D", "2008PhRvD..78l3010G", "2009IJMPD..18.1731M", "2009JCAP...02..034G", "2009JCAP...06..019W", "2009PhRvD..79h3014B", "2009PhRvD..80h4044T", "2009PhRvD..80l3528L", "2010JCAP...08..021B", "2010LRR....13....3D", "2010PhRvD..81j4020F", "2010PhRvD..81l7301L", "2011JCAP...03..002Z", "2011PhR...505...59N", "2012PhRvL.108c1101B", "2012SCPMA..55.2244Z", "2013CQGra..30a5008B", "2013PhRvD..87b3508H", "2013PhRvD..87l3529B", "2013PhRvD..88d4050A", "2013PhRvD..88h4032X", "2014A&A...571A..16P", "2014JCAP...03..046D", "2014JCAP...08..042P", "2014MNRAS.443.1065B" ]
[ "10.1140/epjc/s10052-014-2964-6", "10.48550/arXiv.1403.0655" ]
1403
1403.0655_arXiv.txt
The accelerating expansion of present Universe was discovered by type Ia supernovae \cite{Perlmutter:1998np,Riess:1998cb}. Up to now, the standard $\Lambda$CDM model in the framework of general relativity (GR) is able to explain the present cosmic acceleration within observational errors. However how to explain the tiny value of the cosmological constant compared to the known physical scales is still a big challenge. Modified gravity (MG), for example the $f(R)$ gravity, provides a geometrical origin to the present cosmic acceleration. The basic idea of MG dark energy is that gravity is modified on the cosmological scales when the Ricci scalar $R$ is of order of today's Ricci scalar $R_0$, while GR is recovered in the region of $R\gg R_0$. However it is quite non-trivial to construct a viable $f(R)$ dark energy model which is consistent with both cosmological and local gravity constraints. See some typical viable $f(R)$ dark energy models in \cite{Amendola:2006we,Hu:2007nk,Appleby:2007vb,Starobinsky:2007hu,Tsujikawa:2007xu,Cognola:2007zu,Linder:2009jz}. It is useful to introduce the effective equation of state parameter $w=p_{\rm de}/\rho_{\rm de}$ to describe the difference between Friedmann-Robertson-Walker (FRW) background evolutions of MG and the standard $\Lambda$CDM model, where the effective pressure $p_{\rm de}$ and energy density $\rho_{\rm de}$ are determined by using the Einsteinian representation of gravitational field equations. On the other hand, since the gravity in MG is different from GR, the evolution of the matter density perturbation $\delta_m\equiv \delta\rho_m/\rho_m$ provides a crucial tool to distinguish MG dark energy model from dark energy model in GR, in particular the standard $\Lambda$CDM model. For simplicity, the growth rate $f_g$ of matter density perturbation can be parametrized by, \cite{Linder:2007hg}, \m f_g\equiv {d\ln \delta_m\over d\ln a}\equiv \Omega_m(z)^{\gamma(z)}, \label{dgamma} \n where $a$ is the scale factor, $\Omega_m(z)$ is the density parameter for dust-like matter at redshift $z$, and $\gamma(z)$ is the so-called growth index. In $\Lambda$CDM model in GR, $w=-1$ and \m \gamma\simeq 6/11, \n \cite{Wang:1998gt,Linder:2005in}. Generically the effect on the matter density perturbation in MG is encoded in the effective Newton coupling constant $G_{\rm eff}$. For simplicity, we introduce a new quantity $g\equiv G_{\rm eff}/G$ to measure the difference between MG and GR. In general, $w$ is time dependent and $g$ is time and scale dependent in MG, and then the growth index $\gamma$ is expected to be time and scale dependent. During deep matter dominant era GR is recovered, while the gravity is modified in the low redshift era when the cosmic acceleration occurs. One can expect that the evolutions of both FRW background and matter density perturbation in MG are too complicated to be solved analytically from the deep matter dominant era to accelerating era. In this paper we focus on the growth of matter density perturbation in the $f(R)$ dark energy model. We suppose that $g$ is parametrized as follows \m g=g_0+g_1(1-\Omega_m), \label{paramg} \n where $g_0$ and $g_1$ are two constants. Here $g_1$ is used to characterize the time-evolution of $g$. Note that both $g_0=g_0(k)$ and $g_1=g_1(k)$ are scale dependent generically. In the deep matter dominant era $(\Omega_m\rightarrow 1)$, GR should be recovered and then $g\rightarrow 1$. But $g$ can deviate from one at low redshift. This parameterization can cover many viable $f(R)$ dark energy models at low redshift. Based on such a parameterization, we analytically solve the equation of motion of $\delta_m$ and work out an analytic formula of the growth index. Furthermore, we find that $g$ can be directly figured out by comparing the observed growth rate $f_g$ to the prediction of $f_g$ in GR. This paper will be organized as follows. In Sec.~2 we briefly review the $f(R)$ dark energy model. In Sec.~3 we analytically calculate the growth index for $f(R)$ dark energy model. Summary and discussion are given in Sec.~4.
To summarize, we analytically calculate the growth index in the $f(R)$ dark energy model. Actually our results are applicable for more general MG dark energy models, for example $f(T)$ dark energy model \cite{Linder:2010py,Bamba:2010ws,Zheng:2010am,Zhang:2012jsa}, as long as the effect on the growth of matter density perturbation from MG is encoded in $g=G_{\rm eff}/G$. As we know, there are two key parameters for MG dark energy model, namely $w$ and $G_{\rm eff}$ (or equivalently $g$). The former parameter determines the expansion history of our Universe, and the latter parameter tells us how the matter density perturbation grows up. Adopting the analytic formula, we find a simple relation between $g$ and the growth rate in Eq.~(\ref{ddeltaa}), and then we propose that $g$ can be directly figured out by comparing the observed growth rate $f_g$ to the prediction of $f_g$ in GR. In the literatures, one would like to use $f_g\sigma_8(z)$ to characterize the growth of matter density perturbation. In this case one can also use our analytic formula to calculate $f_g\sigma_8(z)$ and then fit out $g_0$ and $A$ from the data. Recently the anisotropic clustering of the Baryon Oscillation Spectroscopic Survey (BOSS) CMASS Data Release 11 (DR11) sample was analyzed. The combination of Planck and CMASS implies $\gamma=0.772_{-0.097}^{+0.124}$ and a similar result $\gamma=0.76\pm 0.11$ is obtained when replacing Planck with WMAP9 in \cite{oai:arXiv.org:1312.4611}. Both results deviate from the prediction of $\Lambda$CDM in GR at more than $2\sigma$ level. The large value of $\gamma$ may come from the the large value of $\sigma_8$ from Planck, or it is just a statistical fluctuation. Considering $f_g\sigma_8(z=0.57)=0.419\pm 0.044$ from BOSS CMASS DR11, we obtain $g_0\simeq 0.73$ in the reference $\Lambda$CDM model ($\Omega_m^0=0.315$ and $\sigma_8=0.829$) from Planck \cite{Ade:2013zuv}. A careful data fitting will be done in the near future \cite{huang2014}. In a word, if such a deviation is confirmed in the future, we really need to modify the gravity. Finally see some other aspects on $f(R)$ dark energy model in \cite{DeFelice:2010aj,Nojiri:2010wj,Polarski:2007rr,Wei:2008ig,Gong:2008fh,Gannouji:2008wt,Bamba:2008hq,Wu:2009zy,Biswas:2011ar,Bamba:2012qi,He:2012rf,Basilakos:2013nfa,Abebe:2013zua,Xu:2013tsa,Dossett:2014oia,Pouri:2014nta} etc. \noindent {\bf Acknowledgments} This work was initiated during High1-2014 KIAS-NCTS joint workshop on particle physics, string theory and cosmology. Q.-G.H. is supported by the project of Knowledge Innovation Program of Chinese Academy of Science and grants from NSFC (grant NO. 10821504, 11322545 and 11335012).
14
3
1403.0655
We derive the analytic formula of the growth index for the dark energy model where the effect on the growth of matter density perturbation from modified gravity (MG) is encoded in the effective Newton coupling constant in MG (or equivalently ). Based on the analytic formula, we propose that the parameter can be directly found by comparing the observed growth rate to the prediction of in general relativity.
false
[ "MG", "Newton", "general relativity", "matter density perturbation", "modified gravity", "the observed growth rate", "the dark energy model", "the growth index", "the growth", "the analytic formula", "the effect", "the prediction", "the parameter", "We", "we" ]
11.07646
1.306679
-1
437704
[ "Schilke, P.", "Neufeld, D. A.", "Müller, H. S. P.", "Comito, C.", "Bergin, E. A.", "Lis, D. C.", "Gerin, M.", "Black, J. H.", "Wolfire, M.", "Indriolo, N.", "Pearson, J. C.", "Menten, K. M.", "Winkel, B.", "Sánchez-Monge, Á.", "Möller, T.", "Godard, B.", "Falgarone, E." ]
2014A&A...566A..29S
[ "Ubiquitous argonium (ArH<SUP>+</SUP>) in the diffuse interstellar medium: A molecular tracer of almost purely atomic gas" ]
106
[ "I. Physikalisches Institut der Universität zu Köln, Zülpicher Str. 77, 50937, Köln, Germany", "The Johns Hopkins University, Baltimore, MD, 21218, USA", "I. Physikalisches Institut der Universität zu Köln, Zülpicher Str. 77, 50937, Köln, Germany", "I. Physikalisches Institut der Universität zu Köln, Zülpicher Str. 77, 50937, Köln, Germany", "Department of Astronomy, The University of Michigan, 500 Church Street, Ann Arbor, MI, 48109-1042, USA", "California Institute of Technology, Pasadena, CA, 91125, USA; Sorbonne Universités, Université Pierre et Marie Curie, Paris 6, CNRS, Observatoire de Paris, UMR 8112 LERMA, Paris, France", "LERMA, CNRS UMR 8112, Observatoire de Paris &amp; École Normale Supérieure, 24 rue Lhomond, 75005, Paris, France", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Astronomy Department, University of Maryland, College Park, MD, 20742, USA", "The Johns Hopkins University, Baltimore, MD, 21218, USA", "Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA, 91109, USA", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany", "I. Physikalisches Institut der Universität zu Köln, Zülpicher Str. 77, 50937, Köln, Germany", "I. Physikalisches Institut der Universität zu Köln, Zülpicher Str. 77, 50937, Köln, Germany", "LERMA, CNRS UMR 8112, Observatoire de Paris &amp; École Normale Supérieure, 24 rue Lhomond, 75005, Paris, France", "LERMA, CNRS UMR 8112, Observatoire de Paris &amp; École Normale Supérieure, 24 rue Lhomond, 75005, Paris, France" ]
[ "2014A&A...566L...6M", "2014A&A...567A..53K", "2014A&A...569L...5M", "2014ApJ...789....8N", "2014FaDi..168....9V", "2014RSCAd...462030J", "2014sf2a.conf...17D", "2015A&A...577A..49N", "2015A&A...582L...4M", "2015ARA&A..53..199G", "2015ApJ...800...40I", "2015ApJ...801..141O", "2015ApJ...812...75G", "2015ApJ...814..119H", "2015ApJS..217...10M", "2015EAS....75..295B", "2015JPCA..119.4915T", "2015JPCA..119.6528B", "2015MNRAS.446..195T", "2015MNRAS.447..417P", "2015PSST...24d3001M", "2016A&A...593A..56S", "2016A&A...595A.128M", "2016ARA&A..54..181G", "2016ApJ...820L..26B", "2016ApJ...822..115K", "2016ApJ...826..183N", "2016IAUTA..29..137F", "2016JChPh.145t4302F", "2016JChPh.145w1101M", "2016JMoSp.322...29N", "2016JMoSp.327...95E", "2016MNRAS.455.3281H", "2016MolAs...2...18T", "2017A&A...600A...2W", "2017A&A...604A...6S", "2017A&A...606A.109M", "2017ApJ...845..163N", "2017ESC.....1...60F", "2017MNRAS.465.1137B", "2017MNRAS.469..339S", "2017MNRAS.470.2911B", "2017MNRAS.472.4444P", "2017MolAs...8...27T", "2017PCCP...19.5230N", "2018ASSL..451..219A", "2018ASSL..453....3D", "2018ApJ...860..151B", "2018MNRAS.477..802D", "2018MNRAS.478.1502P", "2018MNRAS.479.2415A", "2018PCCP...2025967W", "2018SSRv..214...63S", "2019A&A...631A..86G", "2019ApJ...871..256A", "2019ApJ...872...55J", "2019ApJ...883...54O", "2019ApJ...885..109B", "2019JChPh.150l4305T", "2019JPCA..123.2544M", "2019MolPh.117.1719G", "2019supe.book..313S", "2020A&A...636A..29W", "2020A&A...643A..91J", "2020ApJ...894...37N", "2020ApJ...902..131D", "2020FrCh....8..462G", "2020JPCA..124.7726T", "2020MNRAS.492.1465S", "2020MolPh.11867813B", "2021A&A...651A...9M", "2021FrASS...8...86M", "2021FrASS...8..213M", "2021FrCh....9..187M", "2021JChPh.154e4303K", "2021JChPh.154m4302T", "2021JChPh.155q4306T", "2021MNRAS.504.3797B", "2022A&A...659A.152J", "2022A&A...664A.190G", "2022A&ARv..30....4G", "2022ARA&A..60..247W", "2022AcA....72..233W", "2022ApJ...930..141J", "2022ApJ...934...53S", "2022ESC.....6.1924D", "2022PCCP...2415824A", "2022PSST...31k4012D", "2022arXiv220805756S", "2023A&A...669L..15L", "2023A&A...670A.111K", "2023A&A...673A..89N", "2023Ap&SS.368...76J", "2023ApJ...955L..26K", "2023Atoms..11...87B", "2023CPL...81940443B", "2023ESC.....7..212D", "2023EleSt...5a3001R", "2023Galax..11...86O", "2023JPCA..127.8083K", "2023JPhB...56a5201T", "2023MNRAS.524.1291P", "2023RMxAA..59..327C", "2024A&A...682A.109M", "2024PCCP...26.7377B", "2024arXiv240300917K" ]
[ "astronomy" ]
26
[ "astrochemistry", "line: identification", "molecular processes", "ISM: abundances", "ISM: molecules", "ISM: structure", "Astrophysics - Astrophysics of Galaxies" ]
[ "1965PPS....86..467K", "1976S&T....52..327C", "1978ApJS...36..595D", "1982ApJS...48...95S", "1982JChPh..76.3529V", "1982JChPh..77..693W", "1982PhyS...25..268B", "1983JChPh..79.2093B", "1984JMoSp.106..124J", "1985JMoSp.113..451R", "1987JChPh..87.2442L", "1988JMoSp.127..279F", "1988JMoSp.128..587B", "1989JChPh..91.6142R", "1990IJMSI..98..179B", "1990LIACo..29..291P", "1994ARA&A..32..191W", "1996A&A...305..960C", "1996A&A...315L..27K", "1998JChPh.109.2242M", "1998JPCRD..27..413H", "1998JQSRT..60..883P", "1998Sci...279.1910M", "1999ApJS..125..237D", "1999JMoSp.195..356O", "2001A&A...370L..49M", "2002A&A...395..357C", "2002RvMG...47...21W", "2002RvMP...74.1015W", "2003A&A...398..621L", "2004JMoSt.695....5C", "2004PhRvA..70b2709T", "2005JMoSp.229..170I", "2005JMoSt.742..215M", "2005JPhB...38..693M", "2005JPhB...38L.175M", "2006A&A...454L..13G", "2007PhRvA..75a2502C", "2008A&A...489..115R", "2008AJ....135.1301V", "2009A&A...495..847G", "2009ApJ...700..137R", "2009ApJ...706.1594N", "2010A&A...518A..26W", "2010A&A...518L...1P", "2010A&A...518L...6D", "2010A&A...518L..42V", "2010A&A...518L.110G", "2010A&A...518L.111O", "2010A&A...521L...1W", "2010A&A...521L...9L", "2010A&A...521L..10N", "2010A&A...521L..12S", "2010A&A...521L..14Q", "2010A&A...521L..20B", "2011A&A...525A..77M", "2012A&A...537A..17R", "2012A&A...540A..87G", "2012A&A...542L...1H", "2012A&A...542L...6N", "2012A&A...544A..19D", "2012A&A...544A..22L", "2012ApJ...745...91I", "2012ApJ...749L..17Y", "2012ApJ...751L..37D", "2012ApJ...752..118S", "2012ApJ...754..105H", "2012ApJ...758...83I", "2012ApJ...759L..26S", "2012JMoSp.280..150N", "2012MNRAS.421.2531M", "2013A&A...549A..21M", "2013A&A...550A..25G", "2013ApJ...762...11F", "2013ApJ...764...25J", "2013Sci...342.1343B", "2014A&A...561A.122L", "2014A&A...566A..30R", "2014ApJ...783L...5C", "2014ApJ...785..135L" ]
[ "10.1051/0004-6361/201423727", "10.48550/arXiv.1403.7902" ]
1403
1403.7902_arXiv.txt
Light hydrides of the type ZH$_{\rm n}$ or ZH$_{\rm n} ^+$ are important diagnostics of the chemical and physical conditions in space. Their lower energy rotational transitions occur for the most part at terahertz frequencies (far-infrared wavelengths). This frequency region can be only accessed to a limited extent from the ground, even at elevated sites, because of strong atmospheric absorptions of H$_2$O and, to a lesser degree, O$_2$ and other molecules. The \textit{Herschel} Space Observatory \citep{Herschel_2010} has provided a powerful new probe of the submillimeter and far-infrared spectral regions, which greatly expands upon the capabilities afforded by \new{earlier} missions, such as the Kuiper Airborne Observatory \citep[KAO;][]{KAO_1976}, the Infrared Space Observatory \citep[ISO;][]{ISO_1996}, and others, or from ground with the Caltech Submillimeter Observatory \citep[CSO;][]{CSO} or the Atacama Pathfinder EXperiment \citep[APEX;][]{APEX_2006}. Observations of hydride molecules, in particular of H$_2$O, in interstellar space, but also in solar system objects, were among the important goals of the \textit{Herschel} mission. In fact, the cationic hydrides H$_2$O$^+$ \citep{H2O+_det_2010}, H$_2$Cl$^+$ \citep{H2Cl+_det_2010}, and HCl$^+$ \citep{HCl+_det_2012} were detected with \textit{Herschel} for the first time in the ISM. While the SH radical has its fundamental transition at a frequency that was inaccessible to the high-resolution Heterodyne Instrument for Far-Infrared Astronomy \citep[HIFI;][]{HIFI_1_2010}, it has been detected \citep{SH_det_2012} with the German REceiver At Terahertz frequencies \citep[GREAT;][]{GREAT_2012} onboard the Stratospheric Observatory For Infrared Astronomy \citep[SOFIA;][]{SOFIA-1_2012,SOFIA-2_2013}. OH$^+$ \citep{OH+_det_2010}. and SH$^+$ \citep{SH+_det_2011} were detected with APEX from the ground shortly before \textit{Herschel}, but many additional observations were carried out with HIFI, \citep[e.g.,][]{Godard2012}. Several hydrides, e.g., OH$^+$ and H$_2$O$^+$, were found to be widespread with surprisingly high column densities, not only in Galactic sources \citep{OH_n^+_2010,H2O+_det_2010,Neufeld2010}, but also in external galaxies \citep{det_H2O+_M82_2010,det_OH+_Mrk231_2010,OH_n^+_Arp220_2013}. As both cations react fast with H$_2$ to form H$_2$O$^+$ and H$_3$O$^+$, respectively, it was suspected that these molecules reside in mostly atomic gas, which contains little H$_2$ \citep{OH_n^+_2010}. Detailed model calculations suggest that the abundances of OH$^+$ and H$_2$O$^+$ are particularly high in gas with molecular fraction of around 0.05 to 0.1 \citep{Neufeld2010,OH_n+_chemistry_2012}. The comparatively high column densities observed for these two molecular cations also require cosmic ray ionization rates in the diffuse ISM to be considerably higher than that in the dense ISM \citep{Neufeld2010,OH_n+_chemistry_2012,CRI-rate_W51_2012}. However, \new{evidence for high ionization rates in the diffuse ISM, in the range $10^{-16} - 10^{-15}$ s$^{-1}$, has been presented already earlier to explain the amount of H$_3 ^+$ in the diffuse ISM \citep{H3+_higher-CRI-rate_1998,models_CRI-rate_diff_2003, IndrioloMcCall2012} }. Even higher cosmic ray ionization rates were estimated for active galaxies such as NGC~4418 and Arp~220 \citep[][$> 10^{-13}$ s$^{-1}$]{OH_n^+_Arp220_2013}. Spectral line surveys of the massive and very luminous Galactic Center sources Sagittarius~B2(M) and (N) were carried out across the entire frequency range of HIFI within the guaranteed time key project HEXOS, \citep{HEXOS_2010}. A moderately strong absorption feature was detected toward both sources near 617.5~GHz, but the carrier proved very difficult to assign \citep{Schilke2010, iaudib}. This feature appears at all velocity components associated with diffuse, foreground gas, but is conspicuously absent at velocities related to the sources themselves, suggesting that the carrier resides only in very diffuse gas. The absorption line was detected toward other continuum sources as well during subsequent dedicated observations (within the guaranteed time key project PRISMAS; \citealp{OH_n^+_2010, iaudib}). Very recently, \citet{Barlow2013} observed a line in emission at the same frequency toward the Crab Nebula supernova remnant, which they assigned to the $J = 1 - 0$ transition of argonium $^{36}$ArH$^+$ at 617.525 GHz. In addition, they observed the $J = 2 - 1$ transition at 1234.602~GHz. Here we present evidence that $^{36}$ArH$^+$ is also responsible for the absorption features detected in the HEXOS and PRISMAS spectra.
We \new{confidently assign} the 617.5~GHz line to the carrier $^{36}$ArH$^+$, since features of $^{38}$ArH$^+$ were also detected toward Sgr~B2(M) and (N) with $^{36}$Ar/$^{38}$Ar ratios close to, but probably smaller than in the solar neighborhood. The line surveys cover the frequency of the $J = 1 - 0$ transition of $^{20}$NeH$^+$ and even though Ne is much more abundant in space than Ar, we do not observe neonium absorption. This \new{difference} is in line with expectations based on the much higher ionization potential of Ne. Our chemical calculations show that ArH$^+$ can exist only in low-density gas with a low H$_2$ fraction ($f({\rm H}_2)\approx 10^{-4}-10^{-3}$), and a weak UV field, while an enhanced cosmic ray flux can boost its abundance. OH$^+$ and H$_2$O$^+$ trace gas with a larger H$_2$ fraction of 0.1, and are therefore complementary probes \citep{OH_n^+_2010, Neufeld2010, OH_n+_chemistry_2012}. It is noteworthy, in this context, that the ArH$^+$ and H$_2$O$^+$ column densities are not well-correlated, although one would assume that ArH$^+$ and H$_2$O$^+$ both trace the stratified PDR structures of diffuse clouds, ArH$^+$ the very outer edge, and H$_2$O$^+$ gas deeper in. It appears that this picture is too simplistic. The aforementioned tracers OH$^+$ and H$_2$O$^+$ trace partly molecular gas, while the so-called CO-dark gas, which is predominantly molecular, but does not contain significant abundances of CO, is best traced by HF, CH, H$_2$O, or HCO$^+$ \citep{Qin2010, Sonnentrucker2010, Flagey2013}, but also by [C{\sc ii}] \citep{Langer2014}, which however is not very specific to this component. The careful analysis of column density variations in these tracers promises to \new{disentangle }the distribution of the H$_2$ fraction, providing a direct observational constraint on the poorly known transition of primarily atomic diffuse gas to dense molecular gas traced by CO emission, putting strong constraints upon magnetohydrodynamical simulations for the interstellar gas \citep[e.g.][]{Micic2012, Levrier2012} and thus potentially evolving into a tool to characterize the ISM. Paradoxically, ArH$^+$ actually is a better tracer of almost purely atomic gas than the H{\sc i} line, because with the column density of H we see gas that could be 0.1\%, 1\%, or 50\% molecular, while ArH$^+$ singles out gas which is more than 99.9\% atomic. However, the possibilities of getting more data are limited. While both the 909~GHz OH$^+$ line and the 607~GHz \emph{para}-H$_2$O$^+$ line can be observed under very good weather conditions from very good sites on the ground \citep[see ,e.g.][]{OH+_det_2010}, ArH$^+$, due to its proximity to the 620.7~GHz water line, is extremely difficult even from excellent sites. Receivers covering these frequencies with SOFIA would therefore be highly beneficial. The other possibility to get access to these species are toward redshifted galaxies. There, however, OH$^+$ and H$_2$O$^+$ are often seen in emission, indicating very different excitation conditions. ArH$^+$ has not been found in extragalactic sources yet, but could be a very good tracer of cosmic rays in diffuse gas with little UV penetration.
14
3
1403.7902
<BR /> Aims: We describe the assignment of a previously unidentified interstellar absorption line to ArH<SUP>+</SUP> and discuss its relevance in the context of hydride absorption in diffuse gas with a low H<SUB>2</SUB> fraction. The confidence of the assignment to ArH<SUP>+</SUP> is discussed, and the column densities are determined toward several lines of sight. The results are then discussed in the framework of chemical models, with the aim of explaining the observed column densities. <BR /> Methods: We fitted the spectral lines with multiple velocity components, and determined column densities from the line-to-continuum ratio. The column densities of ArH<SUP>+</SUP> were compared to those of other species, tracing interstellar medium (ISM) components with different H<SUB>2</SUB> abundances. We constructed chemical models that take UV radiation and cosmic ray ionization into account. <BR /> Results: Thanks to the detection of two isotopologues, <SUP>36</SUP>ArH<SUP>+</SUP> and <SUP>38</SUP>ArH<SUP>+</SUP>, we are confident about the carrier assignment to ArH<SUP>+</SUP>. NeH<SUP>+</SUP> is not detected with a limit of [NeH<SUP>+</SUP>]/[ArH<SUP>+</SUP>] ≤ 0.1. The derived column densities agree well with the predictions of chemical models. ArH<SUP>+</SUP> is a unique tracer of gas with a fractional H<SUB>2</SUB> abundance of 10<SUP>-4</SUP> - 10<SUP>-3</SUP> and shows little correlation to H<SUB>2</SUB>O<SUP>+</SUP>, which traces gas with a fractional H<SUB>2</SUB> abundance of ≈0.1. <BR /> Conclusions: A careful analysis of variations in the ArH<SUP>+</SUP>, OH<SUP>+</SUP>, H<SUB>2</SUB>O<SUP>+</SUP>, and HF column densities promises to be a faithful tracer of the distribution of the H<SUB>2</SUB> fractional abundance by providing unique information on a poorly known phase in the cycle of interstellar matter and on its transition from atomic diffuse gas to dense molecular gas traced by CO emission. Abundances of these species put strong observational constraints upon magnetohydrodynamical (MHD)simulations of the interstellar medium, and potentially could evolve into a tool characterizing the ISM. Paradoxically, the ArH<SUP>+</SUP> molecule is a better tracer of almost purely atomic hydrogen gas than Hi itself, since Hi can also be present in gas with a significant molecular content, but ArH<SUP>+</SUP> singles out gas that is &gt;99.9% atomic.
false
[ "atomic diffuse gas", "diffuse gas", "dense molecular gas", "gas", "determined column densities", "interstellar matter", "ArH", "different H<SUB>2</SUB", "several lines", "multiple velocity components", "gt;99.9% atomic", "ISM", "chemical models", "H<SUB>2</SUB", "cosmic ray ionization", "little correlation", "CO emission", "strong observational constraints", "hydride absorption", "Abundances" ]
11.888072
11.222339
-1
437599
[ "Öttl, S.", "Kimeswenger, S.", "Zijlstra, A. A." ]
2014A&A...565A..87O
[ "Ionization structure of multiple-shell planetary nebulae. I. NGC 2438" ]
13
[ "Institute for Astro and Particle Physics, Leopold Franzens Universität Innsbruck, Technikerstrasse 25, 6020, Innsbruck, Austria", "Instituto de Astronomía, Universidad Católica del Norte, 0610 Avenida Angamos, Antofagasta, Chile ; Institute for Astro and Particle Physics, Leopold Franzens Universität Innsbruck, Technikerstrasse 25, 6020, Innsbruck, Austria", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, Manchester, M13 9PL, UK" ]
[ "2015Ap&SS.357...21A", "2015ApJ...812..133S", "2015RMxAA..51..221Z", "2016ApJS..226...22R", "2016MNRAS.455.1459F", "2018A&A...620A..84B", "2018Galax...6...84B", "2018MNRAS.480.1626B", "2020A&A...640A..10W", "2020A&A...644A..63A", "2020MNRAS.493.2238A", "2022A&A...658A..17S", "2022ApJ...939..103R" ]
[ "astronomy" ]
5
[ "planetary nebulae: general", "planetary nebulae: individual: NGC 2438", "stars: AGB and post-AGB", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1977RMxAA...2..181T", "1983ApJ...271..188K", "1983ApJS...51..211A", "1985A&A...142..441O", "1987ApJS...64..529C", "1988A&A...202..253K", "1988ApJ...334..862M", "1989A&AS...79..329Z", "1989ApJ...343..811S", "1989SSRv...51..339K", "1990ApJ...359..392K", "1991A&A...249..223G", "1992A&A...255..477F", "1992ASPC...25..115W", "1992ApJ...392..582B", "1992secg.book.....A", "1993ApJ...409..720T", "1994MNRAS.271..257K", "1995A&A...293..541V", "1995A&A...297..727B", "1996AJ....112..307C", "1996ApJ...459..606S", "1998AJ....115.1693C", "1998PASP..110..761F", "1999A&A...342..823V", "1999A&A...347..169R", "1999ApJ...522..378G", "1999MNRAS.305..190M", "2000A&A...354.1071C", "2001A&A...374..599B", "2001A&A...377L..18S", "2002A&A...382..282C", "2002RMxAC..12..180A", "2003A&A...400..511M", "2003IAUS..209..191R", "2003IAUS..209..511A", "2003MNRAS.340..417C", "2003PASP..115..170S", "2003RMxAC..15...29O", "2004A&A...414..993P", "2004RMxAA..40..193P", "2005AJ....130..172O", "2006AJ....132.1669S", "2006agna.book.....O", "2007PASP..119.1349M", "2008MNRAS.391..399K", "2009ApJ...700.1067M", "2010A&A...511A..53V", "2010A&A...517A..32P", "2010ApJ...720.1738P", "2010PASA...27..187R", "2011apn5.confP..81D", "2012A&A...543A..92N", "2012ASPC..452...99R", "2012AdSpR..50..843D", "2012ApJ...755...53Z", "2012MNRAS.420.1977R", "2012MNRAS.423.3753R", "2012MNRAS.425.1091R", "2013A&A...557A.121G", "2013RMxAA..49..137F", "2014A&A...561A..10P" ]
[ "10.1051/0004-6361/201323205", "10.48550/arXiv.1403.6715" ]
1403
1403.6715_arXiv.txt
{\it Multiple shell planetary nebulae (MSPNe):} \newline Planetary nebulae (PNe) are the ionized ejecta from an asymptotic giant branch (AGB) star. They are a short-lived phenomenon compared to a typical stellar lifetime, and are visible while the now post-AGB star crosses the Hertzsprung-Russell diagram (HRD) towards high temperatures, before entering the white-dwarf cooling track. Most of the luminous material originates from the stellar wind during the last thermal pulses on the AGB. This paper deals with the physical conditions of a special type of PNe: Multiple-shell planetary nebulae (MSPNe), which are surrounded by faint outer shells and/or haloes. In recent times, an increasing number of haloes and multiple shells around PNe have been discovered. Although MSPNe are a familiar phenomenon, most of them are poorly studied and understood. Previous research (e.g., Corradi et al. \cite{corradi_03}, Zhang et al. \cite{zhang12} or Ramos-Larios et al. \cite{IC418}\&\cite{NGC6369}) identified the existence as common feature for nearly-round nebulae. They appear during the evolution around the knee in the HRD and in the early part of the cooling track. The observed MSPN structures are the intricate result of the interaction of hydrodynamic and radiative processes, both during the AGB and the post-AGB phases. A detailed description of the mass loss history and the connection between the stellar winds and the huge extended circumstellar envelopes can be found in, e.g., Bl{\"o}cker (\cite{bloecker_95}) and Decin (\cite{decin}). Due to the faintness of the haloes, up to a factor of $10^3$ below the main nebula in surface brightness, most of them were discovered much later than the nebula itself. Corradi et al. (\cite{corradi_00}), Sch\"{o}nberner \& Steffen (\cite{schoenberner_02}) and Perinotto et al. (\cite{perinotto_04}) used 1D radiative transfer hydrodynamic (RTH) models to calculate the evolution of this kind of nebulae. They modeled the full evolution of the PN starting at the AGB in a sophisticated way. These important models provide a very good representation of the surface brightness of MSPNe at an age of about 10\,000 years. The RTH models result in outgoing shock fronts. At an age of about 10\,000 years (assuming the Bl{\"o}cker (\cite{bloecker_95}) track of an 0.605\,M$_\odot$ central star), the resulting density profile peaks at $n_{\rm H}\gtrapprox$ 500~cm$^{-3}$ (and thus $n_{\rm e}\gtrapprox$ 560~cm$^{-3}$). The central star of the PN (CSPN) has a luminosity of $\approx$250~$L_\odot$ and a temperature of $T_{\rm CSPN} \approx 120$~kK. In those models, the outer limit of the bright nebula is the limit of the hard UV radiation -- a {\sl radiation bounded} optically thick PN. The shell material is nearly recombined ($n_{\rm e} : n_{\rm H} = 1 : 10$), with a temperature of 2\,000~K. Thus the [\ion{O}{i}] $\lambda$6300\AA~ line is predicted to be about 6 times stronger than \ion{H}{$\beta$}. These kind of models are enormously important for the general understanding of the evolution. The layered structure predicted by the models compares well with observations. In this paper, we will investigate the ionization structure of a MSPN, in order to provide further observational constraints for the models. In this work we use the nomenclature introduced by Chu et al. (\cite{chu_87}) and Balick et al. (\cite{balick_92}), but fully defined by Corradi et al. (\cite{corradi_03}), based on the evolutionary RTH models: \begin{itemize} \item{}the main nebula - also called 'rim' by them; \item{}the first thin surrounding structure including its outer weak intensity enhancement is called 'shell'; \item{} the faint outer structures are called 'halo1' and 'halo2'. \end{itemize} \smallskip \noindent{\it NGC 2438:}\newline Our target is a classical MSPN, used as 'benchmark' for the modeling (Corradi et al. \cite{corradi_00}). NGC 2438 shows a bright inner main nebula; the geometry of the nebula is near round and closed. The diameter of the main nebula is about 60\arcsec. The nebula indicates two slightly detached shells and a very faint halo (Fig. \ref{slits.fig}). This faint halo is most visible in the western part of the nebula and seems to have a circular shape. We can find ray-like structures and clumps in the nebula as well. The CSPN is not the bright star near the center, but the fainter one at the center of the nebula. The only observational study of the nebula over a wide optical wavelength range and through the whole main nebula was obtained by Guerrero \& Manchado (\cite{gu_ma_99}), with a single exposure at the ESO 1.5m telescope. They centered the list on the brightest star near the center (not the CSPN), and integrated over wide areas along the slit. Guerrero \& Manchado (\cite{gu_ma_99}) report no detection of the [\ion{O}{i}] line. This might be due to the strength of the telluric airglow line at this position. Further studies of the innermost regions around the CSPN were given by Torres-Peimbert \& Peimbert (\cite{ToPe77}), Kaler (\cite{Ka83}), Kingsburgh \& Barlow (\cite{KB94}) and Kaler et al. (\cite{Ka90}). All of them were focusing on abundances and they find a mild helium and nitrogen overabundance.\newline The spectral investigation of Corradi et al. (\cite{corradi_00}) covered the regions of [\ion{O}{iii}] $\lambda5007$\AA~ and \ion{H}{$\alpha$} + [\ion{N}{ii}] $\lambda$6548\AA+$\lambda$6584\AA~ with a high resolution of $n 70\,000$. They provide good results for the expansion of the nebula. \newline The CSPN was investigated in detail by Rauch et al. (\cite{rauch_99}), using n on-local thermodynamic equilibrium (NLTE) stellar atmosphere models. The results are $\log(g[\rm cgs]) = 6.62\pm0.22$, $T_{\rm CSPN} = 114\pm10\,$kK, $L_{\rm CSPN} = 570\,L_\odot$, and $M_{\rm CSPN} = 0.56\pm0.01\,\,M_\odot$. The helium overabundance of the CSPN is slightly above the one found in the studies of the inner nebula mentioned before. Rauch et al. (\cite{rauch_99}) report that the nebula luminosity is an order of magnitude above the luminosity of the CSPN. We later show (see Sec. \ref{distance.sec}), that this discrepancy was not caused by the model, but by the photometry from literature they used. The low CSPN mass would imply a slow post-AGB evolution. \newline Based on the line ratios of [\ion{O}{iii}] : \ion{He}{ii} : \ion{H}{i} in the main nebula and the line ratio of [\ion{O}{iii}] : \ion{H}{i} in the shell, photoionization studies (Armsdorfer et al. \cite{armsdorfer_02}, \cite{armsdorfer_03}) state that the shell consists of ionized material. The required amount of ionizing UV photons can be obtained by a clumpy structure of the main nebula, allowing UV photons to escape. Such structures are established for some well-studied PNe, e.g., the Helix Nebula (O'Dell et al. \cite{helix_05}; Matsuura et al. \cite{helix_09}) or the Ring Nebula (Speck et al. \cite{ring_03a}; O'Dell et al. \cite{ring_03b}). In Dalnodar \& Kimeswenger (\cite{apn5}), the positions of the spokes of enhanced intensity in the shell of NGC 2438 was shown to be correlated to holes in the main nebula. This supports the concept of a {\sl matter bounded} structure. Recent studies of the MSPN IC~418 (Ramos-Larios et al. \cite{IC418}) and NGC~6369 (Ramos-Larios et al. \cite{NGC6369}) reveal very similar results: {\sl 'Radial filaments emanate outwards from most of the }[\ion{N}{ii}]~{\sl knots'}. \medskip In this work we compiled own spectroscopy and multi-wavelength imaging data sets. Combining more information and data gives us the opportunity to search for the origin of the reported discrepancies and to draw detailed constraints for multi-dimensional RTH studies of MSPNe. We analyze long slit spectra, narrow-band images, a VLA radio map and HST archival data. We investigate the main nebula and the shell by means of a sophisticated spatially resolved CLOUDY model (Ferland et al. \cite{cloudy}, \cite{cloudy13}).
The observations of NGC 2438 allow us to derive an individual distance of $1.9\pm0.2\,$kpc and a foreground extinction of $E_{\rm B-V}=0\fm16\pm0\fm01$. We confirm its non-membership of the open cluster M~46, in the foreground. The large discrepancy of the nebula luminosity and the CSPN luminosity (Rauch et al. \cite{rauch_99}) is completely resolved with these values. The model of the main nebula indicates that the old MSPN is matter bounded. The filling factor in the inner region is lower than that in the outer part of the main nebula. This is very similar to the observational results of the spatially resolved, younger Ring and Helix nebulae. Although using only four lines in the parameter fitting, the photoionization model shows a nearly perfect representation of all observed lines. The analysis of the shock sensitive tracers indicates that shocks do not contribute to the excitation of low ionized atoms like [\ion{N}{ii}] and [\ion{S}{ii}]. This old nebula is dominated by photoionization. The surface brightness distribution of a few bright lines by the RTH models (Corradi et al. \cite{corradi_00}; Perinotto et al. \cite{perinotto_04}) lead to a fair representation of the whole nebula. However, the excitation and temperature throughout the nebula, and beyond, needs the handling of small scale clumps to get a self consistent model. A combination of the sophisticated hydrodynamical calculations in the RTH models, including effects of turbulence and clump formation with photoionization, would be required for a complete view. The determined temperature of the shell and the missing bright [\ion{O}{i}] lines lead to the conclusion that the shell of NGC 2438 is fully ionized. Even the [\ion{O}{iii}] lines indicate a photoionized state. Final confirmation by observations of shock sensitive lines (e.g. [\ion{S}{ii}] $\lambda$6716\AA/$\lambda$6732\AA), to verify possible other excitations, is still missing. More detailed investigations to derive electron temperatures from lower ionized species (e.g., [\ion{N}{ii}]($\lambda$6548\AA+$\lambda$6584\AA)/[\ion{N}{ii}]$\lambda$5755\AA~ or [\ion{O}{ii}]($\lambda$3727\AA+$\lambda$3729\AA)/[\ion{O}{ii}]($\lambda$7320\AA+$\lambda$7330\AA)), and from elements not influenced by depletion in dust grains, like [\ion{Ar}{iii}] $\lambda$5192\AA/($\lambda$7135\AA+$\lambda$7751\AA), are highly desired for the main nebula out to the shell. As recently shown by Pilyugin et al. (\cite{temperature}), these line ratios can achieve very high accuracies. It is also shown that the combination of those lines with \ion{H}{$\beta$} can directly achieve abundances. Furthermore, a detailed investigation of the high excitation features in the inner hot wind bubble and of the wind itself emerging the CSPN is suggested. It is desired to complete the whole image and to give input to multi-dimensional radiative transfer hydrodynamic calculations.
14
3
1403.6715
Context. In recent times an increasing number of extended haloes and multiple shells around planetary nebulae have been discovered. These faint extensions to the main nebula trace the mass-loss history of the star, modified by the subsequent evolution of the nebula. Integrated models predict that some haloes may be recombining, and thus are not in ionization equilibrium. But parameters such as the ionization state and thus the contiguous excitation process are only poorly known. The haloes are very extended, but faint in surface brightness -10<SUP>3</SUP> times lower than the main nebula. The observational limits call for an extremely well studied main nebula, to model the processes in the shells and haloes of one object. NGC 2438 is a perfect candidate to explore the physical characteristics of the halo. <BR /> Aims: The aim is to derive a complete data set of the main nebula. This allows us to derive the physical conditions from photoionization models, such as temperature, density and ionization, and clumping. These models are used to derive whether the halo is in ionization equilibrium. <BR /> Methods: Long-slit spectroscopic data at various positions in the nebula were obtained at the ESO 3.6 m and the SAAO 1.9 m telescope. These data were supplemented by imaging data from the HST archive and from the ESO 3.6 m telescope and by archival VLA observations. The use of diagnostic diagrams draws limits for physical properties in the models. The photoionization code CLOUDY was used to model the nebular properties and to derive a more accurate distance and ionized mass. <BR /> Results: We derive an accurate extinction E<SUB>B - V</SUB> = 0.16, and distance of 1.9 ± 0.2 kpc. This locates the nebula behind the nearby open cluster M46 and rules out membership. The low-excitation species are found to be dominated by clumps. The emission line ratios show no evidence for shocks. The filling factor increases with radius in the nebula. The electron densities in the main nebula are ~250 cm<SUP>-3</SUP>, dropping to ~10-30 in the shell. We find the shell in ionization equilibrium: a significant amount of UV radiation infiltrates the inner nebula. Thus the shell still seems to be ionized. The spatially resolved CLOUDY model supports the hypothesis that photoionization is the dominant process in this nebula, far out into the shell. Previous models predicted that the shell would be recombining, but this is not confirmed by the data. We note that these models used a smaller distance, and therefore different input parameters, than derived by us. <P />Based on observations at ESO, SAAO and VLA, and HST archive.
false
[ "archival VLA observations", "HST archive", "VLA", "multiple shells", "m", "telescope", "ESO", "observations", "ionization equilibrium", "SAAO", "extended haloes", "the main nebula", "haloes", "UV radiation", "data", "the inner nebula", "3.6 m telescope", "ionization", "membership", "photoionization models" ]
10.174423
10.194776
151
482820
[ "Sellwood, J. A.", "Carlberg, R. G." ]
2014ApJ...785..137S
[ "Transient Spirals as Superposed Instabilities" ]
115
[ "Department of Physics and Astronomy, Rutgers University, 136 Frelinghuysen Road, Piscataway, NJ 08854, USA", "Department of Astronomy and Astrophysics, University of Toronto, Toronto, ON M5S 3H4, Canada" ]
[ "2014ApJ...792..122L", "2014ApJ...794..173V", "2014AstL...40..724O", "2014MNRAS.443.2757K", "2014MNRAS.443L...1D", "2014MNRAS.444.3756M", "2014PASA...31...35D", "2014hpcn.conf...54B", "2015AJ....149..116M", "2015ApJ...799..213B", "2015ApJ...800..106W", "2015ApJ...802L..13D", "2015ApJ...806..117F", "2015ApJ...806L..29S", "2015ApJ...809..170P", "2015ApJ...810....9C", "2015CeMDA.123..305G", "2015MNRAS.446.1203P", "2015MNRAS.446.4155K", "2015MNRAS.447.3526K", "2015MNRAS.447.3576D", "2015MNRAS.447L..50B", "2015MNRAS.450..266W", "2015MNRAS.450.2132H", "2015MNRAS.450.2217S", "2015MNRAS.450.4150C", "2015MNRAS.451.2922P", "2015MNRAS.453.1867G", "2015MNRAS.454.2954B", "2016ARA&A..54..667S", "2016ASSL..435..189R", "2016ApJ...822..110K", "2016ApJ...826....2S", "2016ApJ...826L..21S", "2016ApJ...830L..40S", "2016BaltA..25...15B", "2016MNRAS.457.2569M", "2016MNRAS.459..199G", "2016MNRAS.460.2472B", "2016MNRAS.461.2383P", "2016MNRAS.461.2789H", "2016MNRAS.461.3835M", "2016MNRAS.463..459M", "2016MNRAS.463.2210K", "2017ASSL..434...77D", "2017ApJ...835L..18L", "2017ApJ...843..141F", "2017MNRAS.466.4279B", "2017MNRAS.467L..21H", "2017MNRAS.468.3361G", "2017MNRAS.471.2187D", "2017MNRAS.471.4314M", "2017RMxAA..53..257V", "2018A&A...616A..11G", "2018ApJ...852..133S", "2018MNRAS.474.5645P", "2018MNRAS.476.1561D", "2018MNRAS.477.1451F", "2018MNRAS.478..932H", "2018MNRAS.478.1576H", "2018MNRAS.481.2041H", "2018MNRAS.481.3794H", "2019A&A...628A..38S", "2019ApJ...872..205L", "2019ApJ...874..177M", "2019ApJ...876....6M", "2019ApJ...877...64D", "2019ApJ...883...77S", "2019ApJ...884....3S", "2019ApJ...884..173K", "2019MNRAS.484.3154S", "2019MNRAS.484.3198D", "2019MNRAS.485..150D", "2019MNRAS.486.1167B", "2019MNRAS.488.4674V", "2019MNRAS.489..116S", "2019MNRAS.489.5165K", "2019MNRAS.490..478L", "2019MNRAS.490.1026H", "2020ApJ...900..150Y", "2020ApJ...900..186E", "2020IAUS..353...65W", "2020MNRAS.493.2111G", "2020MNRAS.496..767B", "2020MNRAS.496.1610A", "2020MNRAS.497.2442F", "2020MNRAS.498.1159P", "2020MNRAS.499.5623Q", "2020PhRvE.102d2108S", "2020arXiv201114812H", "2021A&A...652A.162C", "2021MNRAS.500.5043S", "2021MNRAS.506.3018S", "2021MNRAS.506.3098D", "2021MNRAS.507.2241L", "2022A&A...658A..50A", "2022A&A...661A..98Y", "2022ARA&A..60...73S", "2022MNRAS.512..366A", "2022MNRAS.513..768C", "2022MNRAS.517.2610S", "2022Univ....8..649F", "2022arXiv220515725D", "2023A&A...680A..85S", "2023ApJ...945....3S", "2023JCAP...08..044S", "2023MNRAS.518.1022S", "2023MNRAS.524...18M", "2023MNRAS.525.4161H", "2023arXiv230804069S", "2023arXiv230908930S", "2023arXiv231004508L", "2024ApJ...966...62M", "2024MNRAS.527.2991C", "2024MNRAS.528.5286H" ]
[ "astronomy" ]
14
[ "galaxies: kinematics and dynamics", "galaxies: spiral", "galaxies: structure", "instabilities", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1088/0004-637X/785/2/137", "10.48550/arXiv.1403.1135" ]
1403
1403.1135_arXiv.txt
The origin of spiral patterns in galaxies still has no fully satisfactory dynamical explanation \citep[see reviews by][]{BT08, Sell13a}. Compelling observational evidence, both photometric \citep{Schw76, GGF95, GPP04, ZCR09} and kinematic \citep{Viss78, Chem06, Shet07}, indicates that spiral patterns are gravitationally driven density waves in the stellar disk. \citet{KN79} and \citet{KKC11} found that the best developed, regular patterns are observed in interacting galaxies and perhaps also those with bars. While the behavior in these cases may not be entirely understood, at least the driving mechanism is clear. The majority of disk galaxies with a significant gas component \citep[\eg][]{Oort62} display less regular patterns, whose origin is not so easily accounted for. The spiral patterns may have some vague rotational symmetry, which is predominantly 2- or 3-fold \citep[][their Table 2]{Davi12}, but it is usually far from perfect as arms bifurcate and/or fade with radius. The ubiquity of the phenomenon, taken together with the fact that simulations of isolated, unbarred, cool collisionless stellar disks (cited below) always manifest similar patterns, argues that there must be a mechanism for self-excitation, which is the question we address here. No galaxy in our hierarchical universe is truly isolated, and infalling subhalos are predicted to bombard the outer halo of every galaxy \citep[\eg][]{Boyl10}. Since tides can excite spiral responses, it is possible some patterns are excited by halo substructure \citep[\eg][]{Dubi08}. But the inner halos of galaxies, where fragile thin disks reside, are quite smooth \citep{Gao11} and even large subhalos, such as that which hosted the Sagittarius dwarf galaxy \citep{Belo06}, can be severely tidally disrupted before perturbing the disk \citep{Purc11}. We argue here that spiral patterns appear so readily as self-excited instabilities that disk responses to diffuse sub-halo perturbations probably give rise to a minority of spirals. Continuously changing recurrent transient patterns have been reported over many years from simulations of isolated, unbarred disk galaxy models \citep[\eg][]{HB74, SC84, Rosk08} and no qualitatively different behavior has emerged as the numerical quality has improved. Claims of long-lived spiral modes have not proven to be reproducible \citep{Sell11}. Spiral activity fades over time as stellar random motions rise due to scattering by the fluctuating spiral patterns, but a reasonable amount of gas infall and dissipation can prolong spiral activity apparently indefinitely \citep[\eg][]{SC84, CF85, Toom90}, which is also consistent with modern galaxy formation simulations \citep[\eg][]{Ager11}. Here, we finally address the issue that was left unexplained in \citet[][hereafter \SC]{SC84}, namely the nature of the spirals that arise in such simulations. In a follow-up study to our original work \citep{Sell89}, we reported that the continuously changing patterns appeared to result from the superposition of a few longer-lived waves, each of which had well-defined frequencies and shapes and lasted for between five and ten full rotations of the pattern. These properties are consistent with them being modes, or standing waves, although they did not last indefinitely. Here we provide stronger evidence and propose a mechanism for this interpretation, using simulations of altogether higher quality than those in our original study.
We have presented evidence that spiral activity in simulations of cool, unbarred, collisionless stellar disks results from a recurrent cycle of transient spiral modes of spiral form (Fig.~\ref{fig.spirals}). We describe them as modes because they start as linear instabilities that grow exponentially even from very low-amplitude (Fig.~\ref{fig.ampl20M}) before saturating and decaying. Growing modes are standing wave oscillations that have positive feedback to cause instability. We argue that the growing wave train reflects off corotation, where it is swing-amplified, and again at some inner radius, where the distribution function has been modified by previous disturbances in the disk. Scattering of particles at resonances causes localized heating over a moderately narrow range of angular momenta (Fig.~\ref{fig.actall}), which introduces abrupt changes to the impedance experienced by traveling waves. Such changes cause partial reflection of a subsequent wave, allowing a standing wave, or unstable mode to develop. Spiral instabilities saturate as a result of the onset of horseshoe orbits that appear as the relative overdensity near corotation approaches $\sim 20$\%, as originally proposed by \citet{SB02}. Once growth is halted by the dispersal of the overdensity at corotation, the wave action stored in the remaining disturbance propagates away from corotation until it is absorbed by wave-particle interactions that cause further localized heating. Thus, the instability cycle is able to repeat. The repeated scattering of particles at different locations leads to a general rise of random motion over the disk that weakens its ability to support further coherent waves, and activity gradually fades on a time-scale of some twenty disk rotations. We have also shown that this timescale is longer in low-mass disks because the multi-arm patterns that are dynamically favored in this case transport angular momentum over a shorter radial distance and, therefore, release less energy into random motions (\S\ref{sec.slowheat}). While we recognize that we have not substantiated all the details, we have presented considerable evidence to support our broad picture. In particular, we find the apparent rapidly changing spirals result from the superposition of a small number of relatively long-lived coherent waves (Fig.~\ref{fig.pspectra}); the phase coherence and large limiting amplitude of these waves are most naturally accounted for by unstable modes. The evolution of each disturbance creates the seeds for a fresh instability, since we find more vigorous growth in simulations that are restarted after scrambling only the azimuthal phases of the particles (Fig.~\ref{fig.ampl20M}). Not only does this result support our picture, but it shows that the activity in the simulations owes nothing to pre-existing density structures, that were erased by scrambling. Fig.~\ref{fig.actall} presents evidence of resonance scattering, which we argue changes the impedance of the disk to traveling waves, thereby creating features that cause partial reflection of the waves, allowing fresh cavity modes to develop. Much of our picture builds on previous work by many authors: the dispersion relation for spiral waves \citep{LS66}, their group velocity \citep{Toom69}, swing-amplification \citep{GLB65, JT66, Toom81}, resonance scattering \citep{LBK72, Mark74}, feedback loops \citep{Mark77, Toom81}, global mode analyses \citep{Zang76, Kaln77, ER98}, and horseshoe orbits at corotation \citep{SB02}. We could not have reached our present level of understanding without all of these contributions, yet our picture is distinct from any previous suggestion. In particular, we argue that the assumption of a smooth distribution function, which many authors regard as the natural starting point, fundamentally discards the spiral baby with the bathwater! Direct observational tests of the generating mechanism for spirals are difficult to devise. Evidence for density waves was summarized in the introduction, but does not help to distinguish between rival theories for their origin. However, analysis of the complete phase-space information for a sample of solar neighborhood stars \citep{Sell10} revealed evidence for a resonant scattering feature, of the kind illustrated in Fig.~\ref{fig.actall}, which supports the picture we present here. Furthermore, the existence of such a feature is inconsistent with other leading theories for the origin of spiral patterns \citep[\eg][]{Bert89, Toom90}.
14
3
1403.1135
We present evidence that recurrent spiral activity, long manifested in simulations of disk galaxies, results from the superposition of a few transient spiral modes. Each mode lasts between 5 and 10 rotations at its corotation radius where its amplitude is greatest. The scattering of stars as each wave decays takes place over narrow ranges of angular momentum, causing abrupt changes to the impedance of the disk to subsequent traveling waves. Partial reflections of waves at these newly created features allows new standing-wave instabilities to appear that saturate and decay in their turn, scattering particles at new locations, creating a recurring cycle. The spiral activity causes the general level of random motion to rise, gradually decreasing the ability of the disk to support further activity unless the disk contains a dissipative gas component from which stars form on near-circular orbits. We also show that this interpretation is consistent with the behavior reported in other recent simulations with low-mass disks.
false
[ "disk galaxies", "subsequent traveling waves", "waves", "new locations", "spiral activity", "further activity", "new standing-wave instabilities", "angular momentum", "other recent simulations", "abrupt changes", "narrow ranges", "low-mass disks", "random motion", "simulations", "particles", "a few transient spiral modes", "stars", "each wave decays", "near-circular orbits", "place" ]
10.11111
6.786199
-1
813545
[ "Joergensen, Jakob", "Sannino, Francesco", "Svendsen, Ole" ]
2014PhRvD..90d3509J
[ "Primordial tensor modes from quantum corrected inflation" ]
38
[ "CP-Origins and The Danish IAS, University of Southern Denmark, Campusvej 55, DK-5230 Odense M, Denmark", "CP-Origins and The Danish IAS, University of Southern Denmark, Campusvej 55, DK-5230 Odense M, Denmark", "CP-Origins and The Danish IAS, University of Southern Denmark, Campusvej 55, DK-5230 Odense M, Denmark" ]
[ "2014GReGr..46.1786N", "2014JCAP...06..045K", "2014JCAP...07..020C", "2014JCAP...08..040D", "2014JHEP...06..080S", "2014PhLB..734...41G", "2014PhRvD..89j3524K", "2014PhRvD..89j3525C", "2014PhRvD..90b3525B", "2014PhRvD..90d3505B", "2014PhRvD..90d7303C", "2014PhRvD..90l4061B", "2014SCPMA..57.1460C", "2015EPJC...75..344B", "2015IJMPD..2441001C", "2015JCAP...03..029B", "2015JHEP...02..050C", "2015NuPhB.892..429C", "2015PhLB..746..242C", "2015PhRvD..91b3509F", "2015PhRvD..91d3509B", "2015PhRvD..91f3509G", "2015PhRvD..91j3521N", "2015PhRvD..91l3527R", "2016Ap&SS.361..210G", "2016PhRvD..93f3528F", "2016PhRvD..94b3521C", "2016PhRvD..94b4009R", "2017IJMPD..2630023C", "2017arXiv170103814B", "2019JPhCS1127a2018B", "2021ApJ...907..107M", "2021PhRvD.104e5030M", "2022EPJC...82..504S", "2022EPJST.231.1325S", "2022ForPh..70f0024Y", "2022IJMPD..3150074C", "2024arXiv240304316C" ]
[ "astronomy", "physics" ]
4
[ "98.80.Cq", "Particle-theory and field-theory models of the early Universe", "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1973PhRvD...7.1888C", "1976PhRvD..13.3333G", "1979JETPL..30..682S", "1979PZETF..30..719S", "1980PhLB...91...99S", "1981JETPL..33..532M", "1981PZETF..33..549M", "1981PhRvD..23..347G", "1982PhLB..108..389L", "1982PhRvL..48.1220A", "1982PhRvL..49.1110G", "2008PhLB..659..703B", "2010PhRvD..82d3502O", "2011JCAP...05..007C", "2012PhRvD..86f3513B", "2012PhRvD..86l5035C", "2014A&A...571A..22P", "2014EPJC...74.3022L", "2014JCAP...06..045K", "2014JCAP...10..035M", "2014PDU.....5...75M", "2014PhRvD..89j3525C", "2014PhRvL.112x1101B", "2014PhRvL.112x1301H", "2014arXiv1405.1390A" ]
[ "10.1103/PhysRevD.90.043509", "10.48550/arXiv.1403.3289" ]
1403
1403.3289_arXiv.txt
\label{nonminsec} The underlying origin of the inflationary paradigm constitutes a prominent problem in cosmology \cite{Starobinsky:1979ty,Starobinsky:1980te,Mukhanov:1981xt,Guth:1980zm,Linde:1981mu,Albrecht:1982wi}. Inflation is traditionally modelled via the introduction of new scalar fields. Many models have been put forward to describe the dynamics of these scalar fields and their interactions with other fields, as it has been recently reviewed in \cite{Martin:2013tda}. On general grounds any renormalizable field theory will recieve quantum corrections to the potential. One can think of the E. Weinberg and Coleman perturbative quantum corrections to the classical scalar potential of any field theory as a simple example of these type of corrections \cite{Coleman:1973jx, Gildener}. We phenomenologically characterize these corrections to the $\phi^4$ theory by introducing a real parameter $\gamma$ as follows: \begin{align} V_{eff} =\lambda \phi^4 \left(\frac{\phi}{\Lambda} \right)^{4 \gamma } \ , \end{align} with $\Lambda$ a given energy scale. Of course, model by model, one can compute the specific potential as in \cite{Okada:2010jf}. Nevertheless we will show that it is possible to provide useful information on a large class of models corresponding to different values of $\gamma$ using this simple approach. For completeness we analyze the cases in which $\phi$ couples both minimally and non-minimally to gravity. We find that for the non-minimally coupled case, the recent results by BICEP2 indicating the presence of primordial tensor modes \cite{Ade:2014xna} constrains $\gamma$ to lie in the region $0.08-0.12$, at the two-sigma confidence level. However, independently on the validity of the BICEP2 results \cite{Audren:2014cea,Mortonson:2014bja}, it is fundamental to know whether quantum corrected potentials can account for nonzero tensor modes. Interestingly, we also discover that for large primordial tensor modes the results are largely independent on the number of e-foldings. Relevant examples of non-minimally coupled models are the Higgs-Inflation model \cite{Bezrukov:2007ep} and the ones in which the inflaton is a composite state \cite{Channuie:2011rq, Bezrukov:2011mv,Channuie:2012bv,Channuie:2013lla}.
14
3
1403.3289
We analyze quantum corrections on the naive ϕ<SUP>4</SUP> inflation. These typically lead to an inflaton potential which carries a noninteger power of the field. We consider both minimal and nonminimal couplings to gravity. For the latter case we also study unitarity of inflaton-inflaton scattering. Finally we confront these theories with the Planck and BICEP2 data. We demonstrate that the presence of nonvanishing primordial tensor modes requires sizable quantum departures from the ϕ<SUP>4</SUP>-inflation model for the nonminimally coupled scenario which we parametrize and quantify. We compare the results with the minimally coupled case and elucidate the main distinctive features.
false
[ "sizable quantum departures", "primordial tensor modes", "quantum corrections", "gravity", "quantum", "the main distinctive features", "inflaton-inflaton scattering", "unitarity", "the ϕ<SUP>4</SUP>-inflation model", "a noninteger power", "the naive ϕ<SUP>4</SUP> inflation", "an inflaton potential", "the field", "the Planck and BICEP2 data", "the nonminimally coupled scenario", "the minimally coupled case", "the latter case", "both minimal and nonminimal couplings", "the Planck and BICEP2", "the presence" ]
11.017626
-0.573876
89
437538
[ "Marshall, J. P.", "Moro-Martín, A.", "Eiroa, C.", "Kennedy, G.", "Mora, A.", "Sibthorpe, B.", "Lestrade, J. -F.", "Maldonado, J.", "Sanz-Forcada, J.", "Wyatt, M. C.", "Matthews, B.", "Horner, J.", "Montesinos, B.", "Bryden, G.", "del Burgo, C.", "Greaves, J. S.", "Ivison, R. J.", "Meeus, G.", "Olofsson, G.", "Pilbratt, G. L.", "White, G. J." ]
2014A&A...565A..15M
[ "Correlations between the stellar, planetary, and debris components of exoplanet systems observed by Herschel" ]
57
[ "Depto. de Física Teórica, Universidad Autónoma de Madrid, Cantoblanco, 28049, Madrid, Spain ; School of Physics, University of New South Wales, Sydney, NSW, 2052, Australia", "Department of Astrophysics, Center for Astrobiology, Ctra. de Ajalvir, km 4, Torrejon de Ardoz, 28850, Madrid, Spain; Space Telescope Science Institute, 3700 San Martin Dr, Baltimore, MD, 21218, USA", "Depto. de Física Teórica, Universidad Autónoma de Madrid, Cantoblanco, 28049, Madrid, Spain", "Institute of Astronomy (IoA), University of Cambridge, Madingley Road, Cambridge, CB3 0HA, UK", "ESA-ESAC Gaia SOC., PO Box 78, 28691, Villanueva de la Cañada, Madrid, Spain", "SRON Netherlands Institute for Space Research, 9747 AD, Groningen, The Netherlands", "Observatoire de Paris, CNRS, 61 Av. de l'Observatoire, 75014, Paris, France", "Depto. de Física Teórica, Universidad Autónoma de Madrid, Cantoblanco, 28049, Madrid, Spain; INAF Observatorio Astronomico di Palermo, Piazza Parlamento 1, 90134, Palermo, Italy", "Department of Astrophysics, , Centre for Astrobiology (CAB, CSIC-INTA), ESAC Campus, PO Box 78, 28691, Villanueva de la Cañada, Madrid, Spain", "Institute of Astronomy (IoA), University of Cambridge, Madingley Road, Cambridge, CB3 0HA, UK", "Herzberg Astronomy &amp; Astrophysics, National Research Council of Canada, 5071 West Saanich Rd, Victoria, BC, V9E 2E7, Canada; University of Victoria, Finnerty Road, Victoria, BC, V8W 3P6, Canada", "School of Physics, University of New South Wales, Sydney, NSW, 2052, Australia; Australian Centre for Astrobiology, University of New South Wales, Sydney, NSW, 2052, Australia; Computational Engineering and Science Research Centre, University of Southern Queensland, Toowoomba, 4350, Queensland, Australia", "Department of Astrophysics, , Centre for Astrobiology (CAB, CSIC-INTA), ESAC Campus, PO Box 78, 28691, Villanueva de la Cañada, Madrid, Spain", "Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA, 91109, USA", "Instituto Nacional de Astrofísica, , Óptica y Electrónica, Luis Enrique Erro 1, Sta. Ma. Tonantzintla, Puebla, Mexico", "SUPA, School of Physics and Astronomy, University of St. Andrews, North Haugh, St. Andrews, KY16 9SS, UK", "UK Astronomy Technology Centre, Royal Observatory Edinburgh, Blackford Hill, Edinburgh, EH9 3HJ, UK; Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh, EH9 3HJ, UK", "Depto. de Física Teórica, Universidad Autónoma de Madrid, Cantoblanco, 28049, Madrid, Spain", "Department of Astronomy, Stockholm University, AlbaNova University Center, Roslagstullsbacken 21, 106 91, Stockholm, Sweden", "ESA Astrophysics &amp; Fundamental Physics Missions Division, ESTEC/SRE-SA, Keplerlaan 1, 2201 AZ, Noordwijk, The Netherlands", "Department of Physical sciences, The Open University, Walton Hall, Milton Keynes, MK7 6AA, UK; Rutherford Appleton Laboratory, Chilton, OX11 0QX, UK" ]
[ "2014A&A...567A.127S", "2014A&A...570A.114M", "2014ApJ...791..114W", "2014MNRAS.445.3315N", "2014sf2a.conf..181F", "2015A&A...579A..20M", "2015AJ....149...86W", "2015Ap&SS.357..103W", "2015ApJ...799..182P", "2015ApJ...801..143M", "2015MNRAS.449.3121K", "2016A&A...593A..51M", "2016ApJ...826..171G", "2016ApJS..225...15C", "2016Icar..264....1F", "2016MNRAS.459.2893M", "2016MNRAS.461.1850F", "2016MNRAS.462.1735W", "2016MNRAS.462.2285C", "2016P&SS..133...47M", "2016PASP..128j2001B", "2016PhDT.......266M", "2016SSRv..205..213M", "2017MNRAS.467..873C", "2017MNRAS.468.4725H", "2017MNRAS.470.3606H", "2017PASA...34....2N", "2018A&A...617A..76C", "2018ARA&A..56..541H", "2018ApJ...869...10M", "2018MNRAS.476.4584K", "2018MNRAS.480.5560B", "2018arXiv180105850C", "2018exha.book.....P", "2018fdpm.book..255M", "2018haex.bookE.165K", "2019A&A...624A...4M", "2019A&A...630A.132S", "2019MNRAS.483..332G", "2019MNRAS.483.3510V", "2019MNRAS.487.3162C", "2019MNRAS.488...37M", "2019MNRAS.488.3588Y", "2019MmSAI..90..543M", "2020MNRAS.495.1943Y", "2020MNRAS.498.1319M", "2020PASP..132h4401T", "2020PASP..132j2001H", "2021AJ....161..271F", "2021MNRAS.503.1276M", "2022A&A...658A..63M", "2022ASSL..466....3R", "2023A&A...669A...3S", "2023A&A...671A.136D", "2023A&A...680A..64D", "2023AJ....165..238M", "2023MNRAS.523.4801W" ]
[ "astronomy" ]
12
[ "infrared: stars", "infrared: planetary systems", "circumstellar matter", "planet-disk interactions", "Astrophysics - Earth and Planetary Astrophysics" ]
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[ "10.1051/0004-6361/201323058", "10.48550/arXiv.1403.6186" ]
1403
1403.6186_arXiv.txt
Circumstellar debris discs around main sequence stars are composed of second generation dust produced by the attrition of larger bodies \citep{bp93} which are remnants of primordial protoplanetary discs \citep{hernandez07}. Debris discs can be detected and analyzed based on their excess infrared emission from the constituent dust particles. Around 16.4$^{+2.8}_{-2.9}$~\% of main sequence Sun-like stars have evidence of circumstellar dust emission at 70~$\mu$m with \textit{Spitzer} \citep{trilling08}. From observations of FGK stars by the \textit{Herschel}\footnote{\textit{Herschel} is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.} DUNES survey an incidence of 20.2~$\pm$~2.0~\% was measured \citep{eiroa13}, whereas the DEBRIS survey measure an incidence of 16.5~$\pm$~2.5~\% (Sibthorpe et al. in prep). Many circumstellar discs around Sun-like stars are seen to have two temperature components, which has been interpreted as arising from two distinct belts at different stellocentric radii \citep{chen09,morales11}. The cool discs are more commonly seen and analogous to the Edgeworth-Kuiper belt (EKB) in our own Solar system \citep{greaves10,vitense12}. The EKB's existence has been inferred from the detection of over a thousand Trans-Neptunian Objects\footnote{1258 as of 29th October 2013, see:\\ http://www.minorplanetcenter.net/iau/lists/TNOs.html} , through ground-based surveys and in-situ dust measurement from Voyager 1 and 2 \citep{gurnett97} and New Horizons \citep{poppe10,han11}, although direct observation of the dust emission from the EKB is confounded by the bright foreground thermal emission from the zodiacal dust in the inner Solar system \citep{backman95}. The less commonly seen warm debris disc Asteroid-belt analogues, which are more difficult to observe around other stars due to the larger flux density contribution from the stellar photosphere at mid-infrared wavelengths compared to the dust excess, have been detected around $\sim$~2~\% of Sun-like stars \citep[3/7 FGK stars with 24~$\mu$m excess and $T_{\rm dust}~>~100~$K from a sample of 184,][]{trilling08}. Exoplanets\footnote{Databases of exoplanet properties are maintained at http://exoplanet.eu and http://exoplanets.org} around Sun-like stars have been identified through radial velocity \citep[e.g. ][]{cnc,aaps,mayor12} or transit surveys e.g. CoRoT \citep{auvergne09}, WASP \citep{wasp}, HAT \citep{hat} and \textit{Kepler} \citep{kepler}. See \citet{perryman11} for a summary of exoplanet detection techniques. The majority of all exoplanet searches have taken place at optical wavelengths, with a sample focus on mature, Sun-like stars as the most suitable candidates for the radial velocity detection technique. A stars are avoided as their atmospheres lack the narrow lines necessary for radial velocity detections through accurate doppler measurements, however these stars are prime candidates for direct imaging surveys, resulting in the detection of several exoplanet systems around debris disc host stars, e.g. Fomalhaut \citep{kalas08}, HR~8799 \citep{marois08}, $\beta$ Pic \citep{lagrange10} and HD~95086 \citep{rameau13}. M stars have likewise been avoided because they exhibit high levels of stellar variability and their emission peaks in the near-infrared, rendering them noisy and faint, although great efforts have been made to overcome these issues due to their sheer number and potential to yield low mass planets through either transit or radial velocity detections \citep[e.g. ][]{reiners10,rodler11,giacobbe12,at12}. This has led to unavoidable bias in the types of stars around which exoplanets are known to exist. Comparative analysis or aggregation of results from radial velocity surveys is further complicated due to both the differences in sensitivity and the variable baseline of observations for the stellar samples, although broad conclusions on e.g. the absence of Jupiter analogues around most nearby stars may be drawn \citep{cumming99,fischer14}. Recent results from an analysis of microlensing surveys suggest that almost all stars may have one or more exoplanets, with low mass planets being much more common than Jupiter mass ones \citep{exoplanets}. On the other hand, long term monitoring from the ground has constrained the likelihood of Jovian planets on long orbits (3--6~AU) to $<$~30\% of the stellar systems surveyed \citep{wittenmyer10}, suggesting that exoplanetary systems very like our own may be rare, a result supported by recent direct imaging searches \citep{janson13}. Given that planets are the end state of the agglomeration of smaller bodies from dust to planetesimals (neglecting as a detail the capture of an envelope from the protoplanetary disc for gas giants), and debris discs are the result of collisional grinding of these planetesimals into dust, one might expect the presence of planets and debris discs to be correlated. This expectation is strengthened by the direct imaging of several exoplanet systems around debris disc host stars, as previously noted, and indirectly by the structural features observed in many debris discs (warps, off-sets, asymmetries). These features have generally been thought to arise from the gravitational perturbation of exoplanets \citep[see reviews by][]{wyatt08,krivov10,moromartin13}, although remnant gas may offer an alternative explantation in some cases \citep[e.g. ][]{lyra13}. Evidence of disc structures, particularly in the sub-mm, has sometimes weakened upon further scrutiny, e.g. the SCUBA detected 'blob' in Vega's disc \citep{holland98,wyatt03,pietu11,hughes12} or the clumps in the disc of HD~107146 \citep{corder09,hughes11}, requiring that caution be exercised when attributing these structures to unseen planetary bodies. Furthermore, no clear correlation has yet been established between the presence of debris discs and the presence of planets \citep{greaves06,moromartin07b,bryden09,kospal09} and larger scale direct imaging surveys have found little evidence of massive planets around debris disc host stars \citep{wahhaj13,janson13}. On the contrary, \cite{maldonado12} identify a trend between the presence of a debris disc and cool Jupiter exoplanet around a star whilst \cite{wyatt12} identified a possible trend between cool dust and low mass planet host stars. Direct measurement of the spatial distribution of debris in other stellar systems will reveal whether our own EKB is common or unusual. The EKB and its interaction with the outer planets is thought to have played a significant role in the development of life on Earth, having supplied a significant fraction of the impactors thought to have been involved in the Late Heavy Bombardment \citep{gomes05}, and continues to provide bodies for the short period comet population through interaction with Jupiter \citep{hj09}. It may also have been involved in the hydration of the Earth, a topic which is still widely debated \citep{obrien06,hmpj09,bond10,hj10,izidoro13}. As such, the nature of exo-EKBs may play a vital role in the determination of the habitability of their host systems \citep[e.g. ][]{hj10}, and it is therefore vital that we are able to judge whether the Solar system's architecture and EKB are the norm, or unusual. The \textit{Herschel} \citep{herschel_ref} Guaranteed Time (GT) debris disc programme and the Open Time Key Programmes DEBRIS\footnote{http://debris.astrosci.ca} \citep[Disc Emission from Bias-free Reconnaissance Infrared Survey; ][]{matthews10} and DUNES\footnote{http://www.mpia-hd.mpg.de/DUNES/} \citep[DUst around NEarby Stars; ][]{eiroa10,eiroa13} have observed nearby Sun-like stars using the PACS \citep[Photodetector Array Camera and Spectrometer,][]{pacs_ref} instrument at far-infrared wavelengths looking for excess emission due to the presence of circumstellar dust discs analogous to the Solar system's EKB \citep{vitense12}. In this work, we have examined all of the stars from the DEBRIS, DUNES and GT programmes currently believed to host exoplanets. In Section 2, we present the observations used in this work, the data reduction process and the stellar physical parameters to be compared to the dust emission. In Section 3, a summary of the assumptions used to fit models to the new \textit{Herschel} photometry are explained and the calculated disc temperatures and masses are shown. In Section 4, a comparison of the observed disc fractional luminosities (or 3-$\sigma$ upper limits in the case of non-detections) from Section 3 and the stellar properties from Section 2 is presented. Finally, in Section 5, we present our conclusions and recapitulate our findings.
We have presented \textit{Herschel} PACS observations of 37 nearby exoplanet systems from the DUNES and DEBRIS samples aimed at searching for exo-EKB analogues. Excess emission attributable to the presence of a circumstellar debris disc was observed around ten of these stars; for the remaining stars we found no evidence of significant excess emission, including the non-measurement of cold emission around HD~69830, providing a tighter upper limit to any possible cold dust in that system. We have improved on the upper limits for dust detection for the stars in this sample by a factor of two over previous \textit{Spitzer} observations, constraining the possible flux from cold dust in all observed systems to at worst two orders of magnitude greater that of the EKB and in several cases at levels comparable to that of the EKB. We find incidences of $\sim$~30~\% for cool debris discs around exoplanet host stars from the sample examined here, irrespective of the spectral type. Due to the large uncertainties in this measurement (from the small sample sizes), these values are in fact consistent with the incidence of debris discs measured by DUNES 20.2~$\pm$~2.0~\% \citep{eiroa13}. The incidence of debris is seen to decrease around older stars, again with large uncertainties. We have identified several trends between the stellar metallicity, the presence of a debris disc and the mass of the most massive exoplanet around the star. We found that low metallicity stars are more likely to host low mass planets, low metallicity stars are also more likely to have a detectable debris disc and that low mass planets are more likely to be associated with a detectable debris disc. This is consistent with what would be expected from the core accretion planet formation model. Combining these trends we develop a picture for these systems in which the gas is stripped from the protoplanetary disc too quickly for Jovian mass planets to form and the resulting low-mass planets can't scatter planetesimals as strongly as more massive planets if they form in-situ, or as they migrate through a disc from a larger initial semi-major axis \citep{maldonado12,wyatt12}. Furthermore, we find no significant evidence for a trend relating the eccentricity of the innermost planet with the fractional luminosity, suggesting that the known exoplanets in these systems have little influence on the visible presence of dust, which is unsurprising as the two components are well separated. We also find no evidence to support the proposed trend between cold Jupiters and lower dust luminosities proposed in \cite{maldonado12}, though this is unsurprising since the sample analysed here is only a sub-set of those from Maldonado et al's work and the newly discovered \textit{Herschel} debris discs have been found exclusively around low mass planet host stars. As an extention to this analysis a companion paper to this work is in preparation, comparing the exoplanets sample presented here to unbiased control samples of stars with exoplanets, or a debris disc, or without either. This future paper will look for differences in the incidence of dusty debris between these different sub-samples. In future work, observations from the \textit{Herschel} SKARPS Open Time Programme \citep[Search for Kuiper Belts Around Radial-Velocity Planet Stars;][]{kennedy13} will be used to increase the number of radial velocity planet host stars for which far-infrared fluxes are available, clarifying, and hopefully supporting, the trends seen between dust, planet and stellar properties in this, and earlier, works.
14
3
1403.6186
Context. Stars form surrounded by gas- and dust-rich protoplanetary discs. Generally, these discs dissipate over a few (3-10) Myr, leaving a faint tenuous debris disc composed of second-generation dust produced by the attrition of larger bodies formed in the protoplanetary disc. Giant planets detected in radial velocity and transit surveys of main-sequence stars also form within the protoplanetary disc, whilst super-Earths now detectable may form once the gas has dissipated. Our own solar system, with its eight planets and two debris belts, is a prime example of an end state of this process. <BR /> Aims: The Herschel DEBRIS, DUNES, and GT programmes observed 37 exoplanet host stars within 25 pc at 70, 100, and 160 μm with the sensitivity to detect far-infrared excess emission at flux density levels only an order of magnitude greater than that of the solar system's Edgeworth-Kuiper belt. Here we present an analysis of that sample, using it to more accurately determine the (possible) level of dust emission from these exoplanet host stars and thereafter determine the links between the various components of these exoplanetary systems through statistical analysis. <BR /> Methods: We have fitted the flux densities measured from recent Herschel observations with a simple two parameter (T<SUB>d</SUB>, L<SUB>IR</SUB>/L<SUB>⋆</SUB>) black-body model (or to the 3σ upper limits at 100 μm). From this uniform approach we calculated the fractional luminosity, radial extent and dust temperature. We then plotted the calculated dust luminosity or upper limits against the stellar properties, e.g. effective temperature, metallicity, and age, and identified correlations between these parameters. <BR /> Results: A total of eleven debris discs are identified around the 37 stars in the sample. An incidence of ten cool debris discs around the Sun-like exoplanet host stars (29 ± 9%) is consistent with the detection rate found by DUNES (20.2 ± 2.0%). For the debris disc systems, the dust temperatures range from 20 to 80 K, and fractional luminosities (L<SUB>IR</SUB>/L<SUB>⋆</SUB>) between 2.4 ×10<SUP>-6</SUP> and 4.1 ×10<SUP>-4</SUP>. In the case of non-detections, we calculated typical 3σ upper limits to the dust fractional luminosities of a few ×10<SUP>-6</SUP>. <BR /> Conclusions: We recover the previously identified correlation between stellar metallicity and hot-Jupiter planets in our data set. We find a correlation between the increased presence of dust, lower planet masses, and lower stellar metallicities. This confirms the recently identified correlation between cold debris discs and low-mass planets in the context of planet formation by core accretion. <P />Tables 2-4 are available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201323058/olm">http://www.aanda.org</A>
false
[ "dust temperature", "cold debris discs", "dust emission", "dust", "fractional luminosities", "lower planet masses", "upper limits", "Stars form", "planet formation", "flux density levels", "Giant planets", "a faint tenuous debris disc", "the debris disc systems", "larger bodies", "the dust fractional luminosities", "statistical analysis", "core accretion", "gas- and dust-rich protoplanetary discs", "ten cool debris discs", "the calculated dust luminosity" ]
8.994845
13.195257
-1
484071
[ "Bond, Nicholas A.", "Gardner, Jonathan P.", "de Mello, Duilia F.", "Teplitz, Harry I.", "Rafelski, Marc", "Koekemoer, Anton M.", "Coe, Dan", "Grogin, Norman", "Gawiser, Eric", "Ravindranath, Swara", "Scarlata, Claudia" ]
2014ApJ...791...18B
[ "The Rest-frame Ultraviolet Structure of 0.5 &lt; z &lt; 1.5 Galaxies" ]
12
[ "Cosmology Laboratory (Code 665), NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Cosmology Laboratory (Code 665), NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Physics Department, The Catholic University of America, Washington, DC 20064 USA", "IPAC, California Institute of Technology, Pasadena, CA 91125, USA", "IPAC, California Institute of Technology, Pasadena, CA 91125, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Department of Physics and Astronomy, Rutgers University, Piscataway, NJ 08854, USA", "Inter-University Centre for Astronomy and Astrophysics, Pune, India", "Minnesota Institute for Astrophysics, School of Physics and Astronomy, University of Minnesota, Minneapolis, MN 55455, USA" ]
[ "2014MNRAS.445..175G", "2015A&A...576A..51G", "2015AJ....150...31R", "2015ApJ...814...97S", "2016ApJ...817...79H", "2016arXiv161001163H", "2017ApJ...836..136F", "2017ApJ...839...71F", "2018MNRAS.474.3976G", "2019A&A...632A.128B", "2019ApJ...875..152B", "2019arXiv191007043B" ]
[ "astronomy" ]
6
[ "cosmology: observations", "galaxies: formation", "galaxies: high-redshift", "galaxies: structure", "Astrophysics - Astrophysics of Galaxies" ]
[ "1976ApJ...209L...1P", "1985ApJS...59..115K", "1996A&AS..117..393B", "2000ApJ...536..571B", "2000ApJS..131..441K", "2003ApJS..147....1C", "2003SPIE.4854...81F", "2003hstc.conf..337K", "2004A&A...428.1043L", "2004ApJ...600L.107F", "2004ApJS..150....1B", "2004ApJS..155..271S", "2005A&A...434...53V", "2005ApJ...619L...1M", "2005ApJ...619L..35H", "2005ApJ...627..632E", "2006A&A...454..423V", "2006AJ....132..926C", "2006AJ....132.1729B", "2006ApJ...636..592L", "2006ApJ...652..963R", "2007ApJ...659..162T", "2007ApJ...670..237B", "2007ApJS..170..377S", "2008A&A...478...83V", "2008ApJ...685L..27R", "2008ApJ...687...59G", "2009A&A...494..443P", "2009AJ....138..362P", "2009ASPC..419...60S", "2009ApJ...694L.158B", "2009ApJ...697.2057L", "2009ApJ...703.2033R", "2010A&A...512A..12B", "2010PASP..122..439R", "2011ApJS..197...35G", "2011ApJS..197...36K", "2012ApJ...753..114W", "2012ApJS..199....4R", "2012MNRAS.427.1666B", "2013A&A...549A..63K", "2013AJ....146..159T", "2013ApJ...763L...7E", "2013ApJ...779..135W", "2013arXiv1305.2140V", "2015ApJS..221...11K" ]
[ "10.1088/0004-637X/791/1/18", "10.48550/arXiv.1403.7463" ]
1403
1403.7463_arXiv.txt
Observations of galaxies at rest-frame ultraviolet wavelengths ($\lambda \sim1500$~\AA) are important for tracing the evolution of star formation and dust obscuration. Until recently, the study of the structural properties of galaxies in the rest-frame ultraviolet has focused on $z\gtrsim2$, as wavelengths $<3000$~\AA\ and redward of the Lyman Break are easily accessible in the observed-frame optical using the Advanced Camera for Surveys \citep[hereafter, ACS,][]{ACS} on the {\it Hubble Space Telescope}\footnote{Based on observations made with the NASA/ESA Hubble Space Telescope, obtained [from the Data Archive] at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with program number HST-GO-12534.}. Furthermore, the {\it Galaxy Evolution Explorer} \citep{GALEX} allows for the study of galaxy structure at $z\lesssim0.5$ \citep[e.g.,][]{GALEX00,Heckman05,TM07}. With the installation of the Wide Field Camera 3 (hereafter, WFC3), including the UVIS channel, we now have the capability to directly observe the UV emission from hot stars in galaxies at $0.5<z<1.5$, a redshift interval that spans about one third of the history of the Universe. Recently published studies of the morphological properties of $0.5<z<1.5$ galaxies are drawn largely from the Cosmic Assembly Near-IR Deep Extraglactic Legacy Survey \citep[CANDELS,][]{Grogin11,Koekemoer11}, which observed $\sim 0.2$~deg$^2$ of sky in the optical and near-infrared with {\it HST}/ACS and {\it HST}/WFC3, respectively. In one such study, \citet{Wuyts12} performed resolved spectral energy distribution (SED) fitting of 323 star-forming galaxies and found that the majority of recent star formation at $0.5 < z < 1.5$ is occurring in clumps at or near the effective radius. These observations are consistent with theoretical models of gas-rich turbulent disks where clumps are supported by infalling cold streams of gas \citep{Bournaud07,Bournaud09}. There are alternative models involving mergers \citep[e.g.,][]{Robertson08}, which may be important for a subset of galaxies at these redshifts, but fragmented structures in sources with clear rotation curves suggest that this is not the dominant mechanism \citep{Genzel08,FS09,Law09,Shapiro09}. At higher redshifts, rest-frame UV imaging with ACS reveals that most $z > 2$ star-forming galaxies are clumpy, disturbed and disk-like in the rest-frame UV, with only $\sim 30\%$ having light profiles consistent with galactic spheroids \citep[e.g.,][]{Ferguson04,EE05,Lotz06,Ravindranath06,Petty09}. These studies find typical half-light radii of $\sim 2$~kpc at $z \sim 2-3$ and a size evolution that scales approximately as $H^{-1}(z)$. Although the UV wavelength dependence of galaxy structure has not been studied at high redshift, such studies have been carried out on well-resolved galaxies in the local universe. \citet{TM07} found that morphology changes occur as one observes bluer in the UV, with galaxies becoming less concentrated, clumpier and more asymmetric. We can obtain a clearer picture of the young stars in galaxies at $0.5 \lesssim z \lesssim 1.5$ by studying their rest-frame UV emission at $\lambda \sim1000-4000$~\AA. Previous studies of star-forming galaxies in this redshift range were performed without the aid of observed-frame UV imaging \citep[e.g.,][]{Bruce12,Wuyts12,Wuyts13} or with relatively shallow imaging in a single filter \citep{VoyerPhD,Rutkowski12}. In this paper, we use data taken as part of a program (GO 11563: PI Teplitz) to obtain UV imaging of the Hubble Ultra Deep Field \citep[hereafter, UVUDF][]{UVUDF} and study intermediate-redshift galaxy structure in the F336W, F275W, and F225W filters, complementing existing optical and near-IR measurements from the 2012 Hubble Ultra Deep Field \citep[HUDF12,][]{HUDF12} survey. We use AB magnitudes throughout and assume a concordance cosmology with $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_{\rm m}=0.27$, and $\Omega_{\Lambda}=0.73$ \citep{WMAP}. With these values, $1\arcsec = 8.0$~physical~kpc at $z=1$.
\label{sec:discussion} Previous studies of star-forming galaxies at $0.5<z<1.5$ in the CANDELS survey \citep{Wuyts12} revealed that the youngest stellar populations at $0.5 < z < 1.5$ tended to be concentrated in clumps near the effective radius (approximately equivalent to $r_{50}$). In a simple 1 Gyr constant star formation model, they found that stars $<10$~Myr old will contribute $\sim 60$\% of the FUV light and stars $<100$~Myr old will contribute $>90$\%. Therefore, we expect that young star-forming clumps, when present, will tend to set the physical scale on which both FUV and NUV emission are observed and $r_{50}$ should be approximately constant across this rest-frame wavelength range. However, far-UV observations of local Sa-Sb galaxies do reveal differences between the $3000$~\AA\ and $1500$~\AA\ light profiles; in particular, they find that galaxies of type later than S0 exhibit a drop in concentration as one observes further into the FUV \citep{TM07}. They attribute this change primarily to the diminished brightness of bulges at shorter wavelengths. Although the fraction of bulge-dominated galaxies decreases with redshift, we still expect $\sim60$\%\ of our galaxies to be bulge-dominated at $z\sim1$ \citep{Bruce12}. Overall, our results are consistent with these expectations, although we do observe a small decrease in $r_{50}$ ($\sim 5$\%) in the FUV for samples at both $0.5<z<1.5$ and $1.5<z<2.5$. The cause of this decrease is not clear, but it is independent of galaxy size. We also observe a decrease in concentration in the FUV, consistent with results at low redshift. It is only marginal for the sample as a whole ($\Delta C\simeq0.05$, Figure~\ref{fig:ConcDiff}), but the largest galaxies ($r_{3000}>2$~kpc) exhibit a drop of $\Delta C \simeq -0.3$, which we find to be due to a flattening of the central portion of the light profile for $\lambda<1800$~\AA. A few illustrative examples are shown in Figure~\ref{fig:Color3v15}, where we plot the pixel-by-pixel color maps of four $0.5<z<1.5$ galaxies between rest-frame $1500$~\AA\ and $3000$~\AA. We also show NIR cutouts from HUDF12 for comparison. While the majority of the UV emission is blue, with $m_{1500}-m_{3000} \sim 0 - 1$, the region near the rest-optical centroid tend to be redder than the rest of the galaxy. This is likely due to the presence of a bulge or proto-bulge near the center of the galaxy with older stellar populations and/or more dust than the rest of the galaxy. To summarize, we find that a $1500$~\AA\ luminosity-limited sample of galaxies at $0.5<z<1.5$ is both smaller ($\sim 5$\%) and less concentrated ($\Delta C\simeq0.05$) at $1500$~\AA\ compared to $3000$~\AA. While the wavelength dependence of $r_{50}$ is independent $r_{3000}$ at all redshifts studied, the decrease in concentration is more substantial for galaxies with $r\gtrsim2$~kpc at $z\sim1$. At $z\sim2$, concentration is approximately constant across the rest-UV for all but the largest galaxies ($r\gtrsim4$~kpc). While we have painted a broad picture of the structural properties of star-forming galaxies in the FUV, a careful analysis of the spatial, size, and color distribution of star-forming clumps is underway (de Mello et al, in prep) and should provide us with a more detailed picture of the star formation in these galaxies.
14
3
1403.7463
We present the rest-frame UV wavelength dependence of the Petrosian-like half-light radius (r <SUB>50</SUB>), and the concentration parameter for a sample of 198 star-forming galaxies at 0.5 &lt; z &lt; 1.5. We find a ~5% decrease in r <SUB>50</SUB> from 1500 Å to 3000 Å, with half-light radii at 3000 Å ranging from 0.6 kpc to 6 kpc. We also find a decrease in concentration of ~0.07 (1.9 &lt; C <SUB>3000</SUB> &lt; 3.9). The lack of a strong relationship between r <SUB>50</SUB> and wavelength is consistent with a model in which clumpy star formation is distributed over length scales comparable to the galaxy's rest-frame optical light. While the wavelength dependence of r <SUB>50</SUB> is independent of size at all redshifts, concentration decreases more sharply in the far-UV (~1500 Å) for large galaxies at z ~ 1. This decrease in concentration is caused by a flattening of the inner ~20% of the light profile in disk-like galaxies, indicating that the central regions have different UV colors than the rest of the galaxy. We interpret this as a bulge component with older stellar populations and/or more dust. The size-dependent decrease in concentration is less dramatic at z ~ 2, suggesting that bulges are less dusty, younger, and/or less massive than the rest of the galaxy at higher redshifts.
false
[ "lt", "large galaxies", "different UV colors", "higher redshifts", "z", "z &lt", "SUB>50</SUB", "r <", "SUB>3000</SUB", "clumpy star formation", "concentration", "0.6 kpc", "6 kpc", "disk-like galaxies", "bulges", "more dust", "length scales", "C", "the rest-frame UV wavelength dependence", "the galaxys rest-frame optical light" ]
12.644428
6.471287
165
871001
[ "Kwitter, K. B.", "Méndez, R. H.", "Peña, M.", "Stanghellini, L.", "Corradi, R. L. M.", "De Marco, O.", "Fang, X.", "Henry, R. B. C.", "Karakas, A. I.", "Liu, X. -W.", "López, J. A.", "Manchado, A.", "Parker, Q. A." ]
2014RMxAA..50..203K
[ "The Present and Future of Planetary Nebula Research. A White Paper by the IAU Planetary Nebula Working Group" ]
32
[ "Department of Astronomy, Williams College, Williamstown, MA, USA.", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-", "-" ]
[ "2015A&A...579A..86W", "2015ASPC..493...83M", "2015ApJ...815...69F", "2015RMxAA..51..221Z", "2015arXiv151204129M", "2016A&A...588A..25M", "2016ApJ...832..125H", "2016ApJS..223....6Y", "2016MNRAS.457.3409A", "2017BAAA...59...43M", "2017NatAs...1E.117J", "2017PASA...34....1D", "2017PASP..129h2001P", "2018A&A...618A...3R", "2018ASPC..517..403S", "2018PhDT.......195L", "2018arXiv181006685S", "2019JPhB...52l5201M", "2020A&A...634A..47M", "2020A&A...638A.103C", "2020Galax...8...23A", "2020MNRAS.495.1016R", "2020PASP..132c5001L", "2020arXiv201009595D", "2021IJGMM..1850191G", "2022ApJ...927..100K", "2022PASP..134b2001K", "2023A&A...676A...1W", "2023ApJ...945...11R", "2023Atoms..11...26S", "2023BAAA...64..133P", "2023PhDT.........2C" ]
[ "astronomy" ]
53
[ "ISM: abundances", "planetary nebulae: general", "stars: AGB and post-AGB", "stars: evolution", "Astrophysics - Solar and Stellar Astrophysics" ]
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"2012ApJ...753..172S", "2012ApJ...760..107G", "2012ApJ...761...35M", "2012ApJ...761..121M", "2012ApJ...761..172G", "2012ApJS..202...12S", "2012IAUS..283...67R", "2012IAUS..283..107M", "2012IAUS..283..144G", "2012IAUS..283..251M", "2012IAUS..283..263P", "2012IAUS..283..267A", "2012IAUS..283..344D", "2012IAUS..283..398J", "2012IAUS..283..504S", "2012MNRAS.419..854G", "2012MNRAS.421.2616N", "2012MNRAS.422.3516W", "2012MNRAS.423L..35P", "2012MNRAS.424.2055H", "2012MNRAS.427.3016P", "2012PhDT........21S", "2012PhRvA..85f2506H", "2012RMxAA..48....3L", "2012Sci...338..773B", "2013A&A...549A..69M", "2013A&A...550C...2F", "2013A&A...550L...6G", "2013A&A...552A..12S", "2013A&A...557L..11B", "2013ApJ...771..114P", "2013ApJ...772...20B", "2013CEAB...37..391D", "2013MNRAS.428.2118D", "2013MNRAS.428.3443M", "2013MNRAS.429.2791F", "2013MNRAS.429.3025C", "2013MNRAS.430..599S", "2013MNRAS.433..295H", "2013MNRAS.436..604R", "2013PhDT.......537B", "2014A&A...563L..10V", "2014ApJ...787..111A", "2014MNRAS.438.2642R", "2014MNRAS.439.2014T" ]
[ "10.48550/arXiv.1403.2246" ]
1403
1403.2246_arXiv.txt
The field of planetary nebula (PN) research has continued to mature, incorporating a variety of new and ingenious approaches that leverage the value of PN observations to maximize astrophysical insight. PNe are the gaseous relics of the evolution of low- and intermediate-mass stars, and as such are ubiquitous in the Galaxy and beyond. They are probes of stellar evolution, populations, gas dynamics, dust and molecules. They are also extragalactic probes of metallicity and dynamics in spiral and elliptical galaxies and in the intracluster medium. Spectra of PNe are characteristic and easy to identify; their emission lines are very strong and can be used to derive C, N, O abundances that in turn characterize the stellar progenitor by comparison of the spectra with the yields from stellar evolution theory. PNe can evolve in binary systems and their morphology and chemical history reflect this type of evolution. They can survive in the intracluster medium, and are rare probes of this interesting environment. Extragalactic PNe have been studied to disclose a characteristic luminosity function whose high-luminosity cutoff appears invariant (or almost invariant) across PN populations, providing a good standard candle. After decades of searching in a variety of celestial objects, researchers have observed fullerenes in PNe, the first environment of stellar origin where these molecules have been observed. As this brief list reveals, the modern study of PNe is extremely fruitful, with many connections to adjacent fields of research, including stellar structure and evolution, binary stars, stellar populations, radial metallicity gradients in spiral galaxies and their evolution, as well as galaxy rotation, evolution, merging, and cosmology. This White Paper by the IAU PN Working Group, represents a summary of our science activities, and is an attempt to set the stage for future developments in the field. Its aims are to raise interest and to spark discussion within the international PN community, as well as across other interested communities, and finally, to serve as preparation for the next PN Symposium several years hence.
14
3
1403.2246
We present a summary of current research on planetary nebulae and their central stars, and related subjects such as atomic processes in ionized nebulae, AGB and post-AGB evolution. Future advances are discussed that will be essential to substantial improvements in our knowledge in the field.
false
[ "post-AGB evolution", "atomic processes", "related subjects", "AGB", "planetary nebulae", "substantial improvements", "current research", "Future advances", "their central stars", "the field", "our knowledge", "a summary", "We", "that" ]
10.061872
10.005246
151
747125
[ "Sutter, P. M.", "Elahi, Pascal", "Falck, Bridget", "Onions, Julian", "Hamaus, Nico", "Knebe, Alexander", "Srisawat, Chaichalit", "Schneider, Aurel" ]
2014MNRAS.445.1235S
[ "The life and death of cosmic voids" ]
35
[ "Sorbonne Universités, UPMC Univ Paris 06, UMR7095, Institut d'Astrophysique de Paris, F-75014 Paris, France; CNRS, UMR7095, Institut d'Astrophysique de Paris, F-75014 Paris, France; Center for Cosmology and AstroParticle Physics, Ohio State University, Columbus, OH 43210, USA", "Sydney Institute for Astronomy, University of Sydney, Sydney, NSW 2016, Australia", "Institute of Cosmology and Gravitation, University of Portsmouth, Portsmouth PO1 3FX, UK", "School of Physics &amp; Astronomy, University of Nottingham, Nottingham NG7 2RD, UK", "Sorbonne Universités, UPMC Univ Paris 06, UMR7095, Institut d'Astrophysique de Paris, F-75014 Paris, France; CNRS, UMR7095, Institut d'Astrophysique de Paris, F-75014 Paris, France", "Departamento de Física Teórica, Módulo C-15, Facultad de Ciencias, Universidad Autónoma de Madrid, E-28049 Cantoblanco, Madrid, Spain", "Department of Physics &amp; Astronomy, University of Sussex, Brighton BN1 9QH, UK", "Department of Physics &amp; Astronomy, University of Sussex, Brighton BN1 9QH, UK" ]
[ "2014JCAP...12..013H", "2014MNRAS.443.2983S", "2014arXiv1409.8661A", "2015JCAP...11..036H", "2015MNRAS.450.3239F", "2015PhRvD..92h3531P", "2016IAUS..308..493V", "2016IAUS..308..524S", "2016MNRAS.456.4425C", "2016MNRAS.458.4431W", "2016MNRAS.461.4013C", "2016PhRvD..93j3519P", "2017CMPh...2013901N", "2017JCAP...07..014H", "2017MNRAS.465..482N", "2017MNRAS.468.3381A", "2017MNRAS.468.4822L", "2018MNRAS.475.3262F", "2018MNRAS.481.2933H", "2019AcASn..60...34Z", "2019EPJC...79..106D", "2019JCAP...12..040V", "2019PASA...36...28E", "2020JCAP...12..023H", "2020MNRAS.491.4554Z", "2021ApJ...916L..24T", "2021ApJ...920L...2V", "2021MNRAS.503.2804M", "2021MNRAS.504.5021C", "2022A&A...658A..20H", "2022LRR....25....6M", "2022MNRAS.511L..82Y", "2023MNRAS.526.1495P", "2024MNRAS.52711962E", "2024arXiv240315134W" ]
[ "astronomy" ]
4
[ "cosmology: theory", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1979Natur.281..358A", "1987ApJ...313..505W", "1993ApJ...410..458D", "1993MNRAS.263..481V", "1994ApJ...431...20S", "1995ApJ...452...25R", "1996MNRAS.282..641S", "2003MNRAS.344..715G", "2004ApJ...605....1G", "2004MNRAS.350..517S", "2004ogci.conf...58V", "2005MNRAS.364.1105S", "2006MNRAS.366..467F", "2007MNRAS.380..551P", "2007MNRAS.382..860D", "2008MNRAS.386.2101N", "2008MNRAS.387..933C", "2010MNRAS.403.1392L", "2010MNRAS.404L..89A", "2011ApJS..192...18K", "2011IJMPS...1...41V", "2012ApJ...754..109L", "2012ApJ...761...44S", "2012ApJ...761..187S", "2012MNRAS.420.1648P", "2012MNRAS.421..926P", "2012MNRAS.421.3481L", "2012MNRAS.426..440B", "2013MNRAS.428.3409A", "2013MNRAS.431..749C", "2013MNRAS.434.1435C", "2013MNRAS.434.2167J", "2013MNRAS.436..150S", "2013MNRAS.436.3525R", "2013PhRvL.111x1103S", "2013arXiv1310.5067S", "2014A&A...571A..19P", "2014MNRAS.438.3177S", "2014MNRAS.440.1248N", "2014MNRAS.440.2922M", "2014MNRAS.441..646N", "2014MNRAS.442..462S", "2014MNRAS.442.3127S", "2014MNRAS.443.2983S", "2014PhRvL.112d1304H", "2014PhRvL.112y1302H", "2015A&C.....9....1S" ]
[ "10.1093/mnras/stu1845", "10.48550/arXiv.1403.7525" ]
1403
1403.7525_arXiv.txt
Since cosmic voids are, by definition, relatively empty of matter, they offer a unique and pristine laboratory for studying dark energy~\citep{LavauxGuilhem2011,Sutter2012b}, exotic fifth forces~\citep{Li2012, Spolyar2013}, and the early universe~\citep{Goldberg2004}. They also offer a complementary probe of the growth of structure via their size and shape distributions~\citep{Biswas2010,Bos2012,Clampitt2013}. Recently large catalogs of voids identified in galaxy redshift surveys~\citep{Pan2011,Sutter2012a,Nadathur2014,Sutter2013c} have opened the way for statistical and systematic measurements of void properties~\citep{Ceccarelli2013,Sutter2013b}, and their connections to cosmological parameters~\citep{Planck2013b, Melchior2013}. However, given the promising utility of voids, we still lack a detailed understanding of their life cycles. For example, for a given void observed at low redshift, we do not know when it formed, where it formed, whether it grew to its present size via simple expansion or through mergers, nor whether it will continue expanding or eventually collapse. We also do not understand basic statistics about voids over cosmic time: their formation and merger rates, growth rates, and movement. Such understanding of the life cycles of voids will solidify current void-based cosmological analysis and enable future probes. Also, if we are to use voids as cosmological probes we must understand the impact of their dynamics on any primordial cosmological signal. As identified theoretically by~\citet{Sheth2004} and discussed in the review of~\citet{VandeWey2011b}, void evolution appears intimately tied to its environment: smaller voids tend to appear inside larger overdense surroundings, while larger voids are truly anti-correlated with respect to the matter distribution. Thus smaller voids tend to collapse over time, while larger voids continue to expand. The expansion of larger voids causes their interiors to appear as miniature open universes, with lower-density walls, filaments, and halos~\citep{vdW2004, Aragon2010, Neyrinck2013} evolving in a self-similar pattern. The evolutionary and hierarchical behavior has been modeled in the context of the adhesion approximation~\citep{Sahni1994}, simulations~\citep{vandeWey1993}, and excursion set theory~\citep{Sheth2004}. However, void abundances are still difficult to predict with excursion set formalisms alone~\citep{Jennings2013,Sutter2013a}. While there have been several attempts to improve the initial theoretical result of~\citet{Sheth2004}, such as by adjusting the void growth and destruction parameters~\citep{Furlanetto2006,Daloisio2007,Paranjape2012} and rescaling void sizes~\citep{Jennings2013}, there still remains very little correspondence to voids identified with watershed techniques in galaxy surveys~\citep{Sutter2013c}. We may improve excursion set predictions by directly measuring the growth and destruction rate in cosmological simulations. The shapes of voids offer a particularly interesting cosmological probe, whether by their distribution~\citep[e.g.][]{Bos2012, Li2012} or via an application of the Alcock-Paczynski test~\citep{Alcock1979, Ryden1995, LavauxGuilhem2011, Sutter2012b, Sutter2014b}. However, these tests rely on the assumption that the void identified in a galaxy survey corresponds to a physical underdensity in the dark matter. While this is largely an issue of sparsity and galaxy bias~\citep{Sutter2013b}, the watershed technique may spuriously merge voids even in the dark matter. These voids will erroneously appear as larger voids that are not completely empty and thus have suspect shapes. We can use a detailed merger history to identify such suspect voids. Recently~\citet{Hamaus2013} pointed out that for a given tracer population there exists a \emph{compensation scale}, where the void-matter bias is identically zero. Below this scale, voids generally collapse due to their surrounding overdense walls, while above this scale voids tend to continue expanding~\citep{Ceccarelli2013}. However, these results are based on studies of the velocity profiles and clustering statistics at fixed time. Only by tracing the evolution --- and thereby studying the dynamics --- of voids could one accurately examine the properties of voids in relation to such a compensation scale. Finally, the growth and merger rates of voids are potential cosmological probes, analogous to the growth rate of cosmic structure. The nature of modified gravity and fifth forces can leave fingerprints on the evolution of the void population at high redshift, potentially constraining the properties of dark energy. Unfortunately, to date this remains a largely unexplored topic. Most early studies of voids in simulations focused on visual identification and characterization~\citep[e.g.,][]{White1987}. For example, the pioneering works of~\citet{Dubinski1993}, which discussed the process of void merging, and~\citet{vandeWey1993}, which first noted the hierarchical nature of void buildup, were entirely based on visually examining thin slices of $N$-body simulations. More recent and more sophisticated analyses have focused on void interiors~\citep[e.g.,][]{Gottlober2003, Goldberg2004,Aragon2012, Neyrinck2013} or on statistics at a fixed time such as those discussed above. In this work we present a comprehensive study of the formation, subsequent evolution, and destruction of voids. We use techniques adapted from building halo merger trees~\citep{Srisawat2013} to follow individual voids across cosmic time. This approach allows us to measure their formation time, identify when mergers occur, track their movement and growth, and, when it does happen, record their time of collapse. We translate this information into rates and correlate these rates with void size, which can then inform theoretical and observational results. In the following section we review our simulation setup, void finding approach, merger identification technique, and some definitions to be used throughout the work. In Section~\ref{sec:formation} we focus on the formation time of voids, followed by a discussion in Section~\ref{sec:growth} of their growth and merger histories. In Sections~\ref{sec:movement} and~\ref{sec:destruction} we present an analysis of void movement and destruction rate over cosmological time, respectively. Finally, we conclude in Section~\ref{sec:conclusions} with a brief discussion of implications for theoretical modeling of voids and directions for future work.
\label{sec:conclusions} We have performed a comprehensive analysis of the life cycle --- covering formation, mergers, growth, movement, and destruction --- of cosmic voids. We have adapted merger tree codes originally designed to track the evolution of halos to account for the large spatial extents of voids. By applying this technique to a high-resolution $N$-body simulation, we have gained a clear picture of voids as dynamic objects in the cosmic web. Through the use of a watershed void finder, we are able to classify voids according to their position in a hierarchy and use that to identify key epochs and scales in their evolution. The past life of a cosmic void depends intimately on its place in the void hierarchy. Voids near the top of the hierarchy primarily form at a scale factor of $0.3$, when the density contrasts in the cosmic web become high enough to support their identification and the introduction of dark energy shuts off continued structure formation. These higher-level voids suffer only minor mergers and tend to maintain consistent growth rates over cosmic time. In contrast, voids lying deep in the hierarchy continue to form and have a somewhat more violent life due to the lower-density nature of their surroundings, but even most of these voids have only a single line of descent. The location of a void in the hierarchy is more important than its size: two voids of equal volume can have radically different merger histories depending on their amount of substructure. Voids typically grow at slow rates. However, there is a population of small collapsing voids. These voids tend to live in overdense environments near filaments and walls, as initially pointed out by~\citet{vdW2004}. Their overdense surroundings slowly squeeze them as adjacent larger voids expand. This picture is consistent with the theory developed by~\citet{Sheth2004}, the velocity inflow-outflow analysis of~\citet{Ceccarelli2013}, the clustering study of~\citet{Hamaus2013}, and the density profile studies of~\citet{Hamaus2014} and~\citet{Sutter2013a}. Despite being slowly crushed, after $a=0.3$, these voids never get completely destroyed. Instead, they continue to survive as identifiable voids to the present day. Thus the void destruction rate does not play a significant role in the late-time evolution of voids, and can be ignored in theoretical treatments. Additionally, as pointed out by~\citet{Russell2013}, the collapsing process is completely negligible for all but the smallest voids, although this result is in conflict with analyses based on the adhesion approximation~\citep{Sahni1994}. Finally, voids do not move much throughout their lifetimes. Only small voids in the frothy depths of the hierarchy that undergo several mergers appear to have perturbed macrocenters. Even for these most active of voids, they typically only move a few percent of their effective radii. The combination of small box volume and high resolution limits our study to relatively small ($\sim 1-15$~\hmpc) voids, while voids in larger simulations and galaxy surveys are typically much larger. However, recently~\citet{Sutter2013a} were able to show that many void properties scale as a function of sampling density and galaxy bias. Thus, properties and characteristics of voids studied in one population of tracers can, in principle, be immediately translated to voids in another population. Thus the conclusions that we reach in this work are generally applicable to voids discovered in other simulations and galaxies. We have examined the properties of voids defined using a watershed technique. There are, of course, other plausible definitions of voids (see, for example, the comparison work of~\citealt{Colberg2008}). These different algorithms might give different pictures of void histories, especially formation times, since they usually impose density thresholds. Additionally, there are other approaches to defining merger trees. We have noticed that volume correlations based on particles can give some non-intuitive results: voids that appear to occupy similar positions (based on their macrocenters and effective radii) may not necessarily share any particles. The relationships between particle correlation and macrocenter definition should be investigated further. However, we have applied other merger tree algorithms, such as \textsc{MergerTree} and \textsc{JMerge} (both described in~\citealt{Srisawat2013}), and found qualitatively similar results. Overall, voids live far quieter lives than their overdense counterparts, the halos. Whereas up to $20\%$ of halos have suffered a recent major merger, voids experience essentially \emph{no} major mergers throughout their lifetime. Likewise, while subhalos can be stripped of their mass as they pass through a larger parent halo, subvoids continue to be identifiable even when a supervoid forms around them. The implication is that voids are a much more pure cosmological probe; the fundamental cosmological signal imprinted from initial conditions and modified by dark energy is not corrupted by significant dynamics. Thus lower-redshift cosmological probes, such as the Alcock-Paczynski test and void-galaxy cross-correlations, will not be affected by recent spurious mergers in the void population.
14
3
1403.7525
We investigate the formation, growth, merger history, movement, and destruction of cosmic voids detected via the watershed transform code VIDE in a cosmological N-body dark matter Λ cold dark matter simulation. By adapting a method used to construct halo merger trees, we are able to trace individual voids back to their initial appearance and record the merging and evolution of their progenitors at high redshift. For the scales of void sizes captured in our simulation, we find that the void formation rate peaks at scale factor 0.3, which coincides with a growth in the void hierarchy and the emergence of dark energy. Voids of all sizes appear at all scale factors, though the median initial void size decreases with time. When voids become detectable they have nearly their present-day volumes. Almost all voids have relatively stable growth rates and suffer only infrequent minor mergers. Dissolution of a void via merging is very rare. Instead, most voids maintain their distinct identity as annexed subvoids of a larger parent. The smallest voids are collapsing at the present epoch, but void destruction ceases after scale factor 0.3. In addition, voids centres tend to move very little, less than 10<SUP>-2</SUP> of their effective radii per ln a, over their lifetimes. Overall, most voids exhibit little radical dynamical evolution; their quiet lives make them pristine probes of cosmological initial conditions and the imprint of dark energy.
false
[ "cold dark matter simulation", "dark energy", "void destruction", "void sizes", "cosmological initial conditions", "scale factor", "merger history", "halo merger trees", "individual voids", "cosmic voids", "high redshift", "little radical dynamical evolution", "the median initial void size", "Voids", "voids", "most voids", "a cosmological N-body dark matter", "destruction", "growth", "evolution" ]
11.276533
2.682706
156
800084
[ "Gao, Qing", "Gong, Yungui" ]
2014PhLB..734...41G
[ "The challenge for single field inflation with BICEP2 result" ]
50
[ "MOE Key Laboratory of Fundamental Quantities Measurement, School of Physics, Huazhong University of Science and Technology, Wuhan 430074, PR China", "MOE Key Laboratory of Fundamental Quantities Measurement, School of Physics, Huazhong University of Science and Technology, Wuhan 430074, PR China" ]
[ "2014JCAP...07..020C", "2014JCAP...08..040D", "2014JCAP...10..025H", "2014JCAP...10..048G", "2014JCAP...11..003L", "2014JHEP...07..026H", "2014JHEP...10..137M", "2014MPLA...2950197F", "2014PTEP.2014j3C01K", "2014PhLB..737..191H", "2014PhLB..737..374B", "2014PhLB..738..206L", "2014PhLB..738..254F", "2014PhLB..738..412G", "2014PhRvD..90b3525B", "2014PhRvD..90b3537W", "2014PhRvD..90d3505B", "2014PhRvD..90d7303C", "2014PhRvD..90l4061B", "2014SCPMA..57.1442G", "2014SCPMA..57.1449W", "2014SCPMA..57.1460C", "2014arXiv1405.2775Z", "2014arXiv1407.6948H", "2015EPJC...75...55L", "2015EPJC...75..301L", "2015EPJC...75..344B", "2015IJMPA..3045004Y", "2015IJMPD..2441001C", "2015JApA...36..269F", "2015JCAP...01..037L", "2015JCAP...01..040K", "2015NuPhB.892..429C", "2015PhRvD..91d3509B", "2015PhRvD..91e3008H", "2015PhRvD..91f3509G", "2016Ap&SS.361..210G", "2016CQGra..33t5001Y", "2016IJMPA..3150047D", "2016MNRAS.459.4029L", "2017IJMPD..2650005Y", "2017SCPMA..60i0411G", "2018EPJP..133..491G", "2018JCAP...05..005G", "2018Univ....4...15G", "2021ApJ...907..107M", "2021JMPh...12..781A", "2021PhRvD.103d3516A", "2022EPJP..137..208G", "2023InJPh..97.4117S" ]
[ "astronomy", "physics" ]
4
[ "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]
[ "1982PhRvL..48.1220A", "1983PhLB..129..177L", "1990PhRvL..65.3233F", "2005JCAP...07..010B", "2010JCAP...09..007B", "2012JCAP...02..008H", "2014A&A...571A..16P", "2014A&A...571A..22P", "2014JCAP...02..025A", "2014JCAP...07..014C", "2014NuPhB.882..386C", "2014PhLB..733..283H", "2014PhLB..734...13H", "2014PhLB..734...96N", "2014PhLB..734..134A", "2014PhLB..736..305D", "2014PhLB..737...98A", "2014PhRvD..89j1302M", "2014PhRvD..89j3524K", "2014PhRvD..89j3525C", "2014PhRvD..89l3505M", "2014PhRvD..90b3523B", "2014PhRvD..90d3509J", "2014PhRvL.112a1303K", "2014PhRvL.112q1301L", "2014PhRvL.112x1101B", "2014arXiv1403.3919Z", "2014arXiv1403.4592C", "2014arXiv1403.5043H", "2014arXiv1403.5549C", "2015JCAP...03..044F", "2015PhLB..745..118H" ]
[ "10.1016/j.physletb.2014.05.018", "10.48550/arXiv.1403.5716" ]
1403
1403.5716_arXiv.txt
The detection of the primordial B-mode power spectrum by the BICEP2 collaboration confirms the existence of primordial gravitational wave, and the observed B-mode power spectrum gives the constraint on the tensor-to-scalar ratio with $r=0.20^{+0.07}_{-0.05}$ at $1\sigma$ level for the lensed-$\Lambda$CDM model \cite{Ade:2014xna}. Furthermore, $r=0$ is disfavored at $7.0\sigma$ level. The new constraints on $r$ and the spectral index $n_s$ exclude a wide class of inflationary models. For the inflation model with non-minimal coupling with gravity \cite{Kallosh:2013tua}, a universal attractor at strong coupling was found with $n_s=1-2/N$ and $r=12/N^2$. This model is inconsistent with the BICEP2 result $r\gtrsim 0.1$ at $2\sigma$ level because the BICEP2 constraint on $r$ requires the number of e-folds $N=\sqrt{12/r}\lesssim \sqrt{120}\approx 11$ which is not enough to solve the horizon problem. If we require $N=50$, then $r=0.0048$, so the model is excluded by the BICEP2 result. For the small-field inflation like the hilltop inflation with the potential $V(\phi)=V_0[1-(\phi/\mu)^p]$ \cite{Albrecht:1982wi,Boubekeur:2005zm}, $r\sim 0$, so the model is excluded by the BICEP2 result. Without the running of the spectral index, the combination of {\em Planck}+WP+highL data gives $n_s=0.9600\pm 0.0072$ and $r_{0.002}<0.0457$ at the 68\% confidence level for the $\Lambda$CDM model \cite{Ade:2013zuv,Ade:2013uln} which is in tension with the BICEP2 result. When the running of the spectral index is included in the data fitting, the same combination gives $n_s=0.957\pm 0.015$, $n_s'=dn_s/d\ln k=-0.022^{+0.020}_{-0.021}$ and $r_{0.002}<0.263$ at the 95\% confidence level \cite{Ade:2013zuv,Ade:2013uln}. To give a consistent constraint on $r$ for the combination of {\em Planck}+WP+highL data and the BICEP2 data, we require a running of the spectral index $n_s'<-0.002$ at the 95\% confidence level. For the single field inflation, the spectral index $n_s$ for the scalar perturbation deviates from the Harrison-Zel'dovich value of $1$ in the order of $10^{-2}$, so $n_s'$ is in the order of $10^{-3}$. The explanation of large $r$ and $n_s'$ is a challenge to single field inflation. In light of the BICEP2 data, several attempts were proposed to explain the large value of $r$ \cite{Lizarraga:2014eaa,Harigaya:2014qza,Contaldi:2014zua,Collins:2014yua,Byrnes:2014xua,Anchordoqui:2014uua,Harigaya:2014sua, Nakayama:2014koa,Zhao:2014rna, Cook:2014dga,Kobayashi:2014jga,Miranda:2014wga,Masina:2014yga,Hamada:2014iga,Hertzberg:2014aha, Dent:2014rga,Joergensen:2014rya,Freese:2014nla,Ashoorioon:2014nta,Ashoorioon:2013eia, Choudhury:2014kma,Choudhury:2013iaa,Hotchkiss:2011gz,BenDayan:2009kv}. In this Letter, we use the chaotic and natural inflation models to explain the challenge.
For a single inflaton field with slow-roll, the tensor-to-scalar ratio $r\approx 16\epsilon$ which is linear with the slow-roll parameter $\epsilon$, but the running of the spectral index $n_s'$ depends on the second order slow-roll parameters, so $n_s'$ is at most in the order of $10^{-3}$. The BICEP2 and the {\em Planck} data constrain $n_s'=-0.0221^{+0.011}_{-0.0099}$ and $r=0.20^{+0.07}_{-0.05}$ at the $1\sigma$ confidence level. Both the chaotic and natural inflation are inconsistent with the observation at the $1\sigma$ level. The chaotic inflation with $2<p<3$ and the natural inflation with $f\gtrsim 10M_{pl}$ are marginally consistent with the observation at the 95\% confidence level. In conclusion, it is a challenge to simultaneously explain $r$ as large as $0.2$ and $n_s'$ as large as $-0.01$ for single field inflation. Unless the {\em Planck} and the BICEP2 data can be reconciled without large $n_s'$, the challenge to single field inflation remains.
14
3
1403.5716
The detection of B-mode power spectrum by the BICEP2 collaboration constrains the tensor-to-scalar ratio r=0.20-0.05+0.07 for the lensed-ΛCDM model. The consistency of this big value with the Planck results requires a large running of the spectral index. The large values of the tensor-to-scalar ratio and the running of the spectral index put a challenge to single field inflation. For the chaotic inflation, the larger the value of the tensor-to-scalar ratio is, the smaller the value of the running of the spectral index is. For the natural inflation, the absolute value of the running of the spectral index has an upper limit.
false
[ "single field inflation", "scalar", "the spectral index", "a large running", "an upper limit", "The large values", "r=0.20", "the absolute value", "BICEP2", "the chaotic inflation", "the natural inflation", "this big value", "the running", "the lensed-ΛCDM model", "the larger the value", "the smaller the value", "the BICEP2 collaboration", "B-mode power", "the Planck results", "the tensor-to-scalar ratio" ]
11.432069
-0.637945
89
643722
[ "Hakopian, Susanna" ]
2014IAUS..304...36H
[ "Complex Investigation of SBS Galaxies in Seven Selected Fields" ]
3
[ "Byurakan Astrophysical Observatory (BAO), Armenia" ]
[ "2016ASPC..505..183H", "2021Ap.....64....8H", "2022Ap.....65..297H" ]
[ "astronomy" ]
1
[ "Astrophysics - Astrophysics of Galaxies" ]
[ "1967Ap......3...24M", "1968ApJ...151..393S", "1970ApJ...162L.155S", "1974ApJ...192..581K", "1980A&A....87..142H", "1983Ap.....19..354M", "1985ApJS...57..503F", "1987A&AS...70..281B", "1987ApJS...63..295V", "1990ApJS...72..567G", "1993ApJ...404..551O", "1997A&A...319...52V", "1999IAUS..194..162H", "2004Ap.....47..378H", "2006Ap.....49..437H", "2012Ap.....55....1H" ]
[ "10.1017/S1743921314003238", "10.48550/arXiv.1403.0127" ]
1403
1403.0127_arXiv.txt
The Second Byurakan Spectral Sky Survey (SBS) (\cite[Markaryan \& Stepanyan, 1983]{MarkaryanStepanyan83}) was undertaken by Markarian as a direct continuation of the First Byurakan Survey (FBS) (\cite[Markaryan, 1967]{Markaryan67}) to achieve magnitudes fainter than in FBS in search of active objects. Both surveys were carried out with the 1-m Shmidt type telescope of Byurakan observatory with the use of objective prisms of the same size as the entrance port of the telescope. Main differences of conducting the two surveys are seen on Table 1. Besides, it must be noted that the objects in SBS have been selected not only by presence of uv-excess on their low dispersion spectra, as it was done in FBS, but also by the presence of emission lines, as the second selection criterium. More than 1500 objects, which have been included in the common sample of FBS, are more famous as "Markarian galaxies". About 3000 objects selected in the 65 fields, each of 16sq. deg., which compile the SBS area, are distributed in the three samples, two of which (the sample of quasars and stars) consist of starlike objects.\\ In the third one, the sample of galaxies, initially about 1300 extended objects were included. Later, with the help of sources other than the SBS photographic plates, 200 more objects have been added to this list. As homogeneity of the object selection is important for the purposes of our program, we are working with the original list of SBS galaxies. \begin{table} \begin{center} \caption{ Main differences of conducting the two byurakan surveys.} \label{tab1} {\scriptsize \begin{tabular}{|l|c|ccc|}\hline {\bf Survey} & {\bf FBS } & {\bf}& {\bf SBS } & {}\\\hline {Limiting magnitudes (pg)} & $ {\sim 17.5} $ & &${\sim 19.5}$&\\ \hline {Objective prism (angular degree)} & {1.5} & $ {1.5} $& $ { 3} $& $ { 4}$ \\ \hline Photographic plate & Kodak IIa-F & {Kodak IIIa-J } & {Kodak IIIa-J} & {Kodak IIIa-F } \\ \hline\ Hypersensitization of the plates &- &{i n}&{n i t r o g e n}& { v a p o r } \\ \hline The used filter & - & {- } & { GG495} & {RG2 } \\ \hline Exposure time for lim mag (min) &15 & {30-60}& {60-120}& {630-690}\\ \hline Dispersion (\AA/mm)& 2500 (H\begin{tiny}$\beta$\end{tiny}) & 1800 (H\begin{tiny}$\beta$\end{tiny}) & 850 (H\begin{tiny}$\gamma$\end{tiny}) & 1100 (H\begin{tiny}$\alpha$\end{tiny}) \\ \hline Spectral range (nm)& 350-690 & 350-540 & 490-540 & 630-690\\ \hline Total area (square degrees)& $\sim1700$ & $\sim1000$ & $\sim160$ & $\sim800$\\ \hline \end{tabular}} \end{center} \vspace{1mm} \scriptsize \end{table} \begin{figure}[] \begin{center} \includegraphics[width=4.5in]{area.eps} \caption \tiny {Boundaries of selected fields within distribution of SBS galaxies in the whole area of SBS.} \label{fig1} \end{center} \end{figure}
14
3
1403.0127
It is known that the main criterion for the selection of active objects in the First Byurakan, otherwise Markarian survey was the presence of signs of UV-excess in their low-dispersion spectra. Using the presence of emission lines as the second criteria became real during the Second Byurakan survey because of its improved technique. Extended (not stellated) objects, selected with the use of this criterion, made the main part of the separate sample of SBS galaxies. Originally, this sample included 1286 objects, selected in 65 fields of the survey (16 square degree each), to which, with the help of other sources than the survey, there were later added some objects. We studied a subsample of SBS galaxies in seven selected fields (the deepest according to the V/Vmax criterion), including about the third of the whole sample. The first, already completed phase of this program was started with carrying out a follow-up slit spectroscopy of all, about 500 objects, based on observations with long-slit spectrographs with 6m telescope of SAO Russia and 2.6.m telescope of Byurakan. As a result redshifts were determined, as well as spectral classification was made for all of objects, using the scheme adapted to the spectral material. Besides other, obtained data allowed us to estimate the efficiency of used criteria for the selection of galaxies of different classes of starformation and nuclear activity along the full scale of the apparent magnitudes, including close to the limit values (18.5 &lt; m <SUB>pg</SUB> &lt; 19.5), etc. The fact that the total area of seven fields as the total number of objects in them comparable with these values for the survey as a whole, allows us to extrapolate the results to the whole sample of galaxies as an upper estimate. The second stage is to conduct detailed studies of individual galaxies in the first place, the most interesting in terms of morphology. They are based on panoramic spectroscopy obtained from observations at 6 m telescope of Russia and 2.6m telescope of Byurakan carried out with multipupil spectrographs MPFS and VAGR, correspondently. Processing of the data obtained for more than twenty objects are at different stages (see arXiv:1403.0127 for extended version).
false
[ "lt", "active objects", "objects", "SBS galaxies", "used criteria", "individual galaxies", "2.6.m telescope", "SUB", "galaxies", "Byurakan", "multipupil spectrographs MPFS", "extended version", "SAO Russia", "the Second Byurakan survey", "VAGR", "different stages", "spectral classification", "different classes", "morphology", "other sources" ]
14.85153
6.765548
119