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charlieoneill
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12435065
[ "Wolfson, Ira", "Brustein, Ram" ]
2016arXiv160703740W
[ "Small field models with gravitational wave signature supported by CMB data" ]
4
[ "-", "-" ]
[ "2019ForPh..6700037P", "2019PLoSO..1415287W", "2019arXiv190303548K", "2022arXiv220703150W" ]
[ "astronomy" ]
2
[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1993PhLB..302..171S", "1995PhRvD..52.1739K", "1997PhLB..414...18W", "1997PhRvD..56.3207D", "1997PhRvL..78.1861L", "1999PhR...314....1L", "2001PhLB..500....1M", "2001PhLB..510....1S", "2001PhLB..517..243S", "2001PhRvD..64l3514A", "2002PhRvD..65j1301D", "2002PhRvD..65j3508S", "2002PhRvD..66j3511L", "2004JCAP...08..003S", "2005JCAP...07..010B", "2006JCAP...09..010E", "2007PhRvD..75l3519L", "2008JCAP...07..011B", "2009ApJS..180..330K", "2009JCAP...03..013C", "2010JCAP...09..007B", "2011CQGra..28j3001Y", "2012JCAP...02..008H", "2013ApJS..208...19H", "2014A&A...571A..16P", "2014A&A...571A..22P", "2014JCAP...05..035A", "2014PDU.....5...75M", "2014PhRvD..90f3501M", "2014PhRvD..90l3539G", "2014PhRvL.112x1101B", "2014SPIE.9153E..1NA", "2015ApJ...811..126B", "2015CRPhy..16..960V", "2015PhRvL.114j1301B", "2015ist..book.....B", "2016A&A...594A..13P", "2016A&A...594A..20P", "2018LRR....21....2A", "2019PLoSO..1415287W" ]
[ "10.48550/arXiv.1607.03740" ]
1607
1607.03740_arXiv.txt
} Recent years have shown an increase in cosmological observational data, largely due to the Planck mission \cite{Ade:2015xua}, and the searches for primordial gravitational waves (GW) signal in the cosmic microwave background (CMB) by terrestrial experiments such as BICEP2 and the Keck Array \cite{Ade:2014xna,Ade:2015fwj}. Inflation \cite{Starobinsky:1979ty,Guth:1980zm,Linde:1981mu,Albrecht:1982wi} is widely accepted as a probable model for the origin of our universe, one of the hallmarks of which is the production of GW (for example \cite{Starobinsky:1979ty,Rubakov:1982df}). Over the years the sensitivity for detecting GW in the CMB has improved constantly. The constraints on the tensor-to-scalar ratio $r$ were significantly tightened \cite{Komatsu:2008hk,Ade:2014xna,Ade:2015fwj,Hinshaw:2012aka,Ade:2013zuv,Ade:2015xua,Ade:2015tva} and it is expected that a sensitivity level of $r\lesssim 0.03$ to be reached in the near future \cite{Ahmed:2014ixy}. Furthermore one can optimistically expect the next decade to yield measurements of $r\lesssim 0.001$ or better possibly \cite{Amendola:2016saw}. Hand in hand with this recent influx of observations, constant headway is made in the model building front, as some models become less probable, while others gain dominance. We study a class of models that were proposed by Ben-Dayan \& Brustein \cite{BenDayan:2009kv}, which sport, along with the ability to conform to known observable quantities such as the primordial power spectrum (PPS) scalar index ($n_{s}$), and its running ($n_{run}$), the possibility of generating appreciable amplitude of GW signal. This type of models appear in many fundamental physics frameworks, such as effective field theory, supergravity and string theory. Later, a general discussion followed regarding small field models and the possibility of GW generation \cite{Hotchkiss:2011gz,Antusch:2014cpa,Garcia-Bellido:2014wfa}. In these models, high values of $r$ in the CMB are generally associated with a scale dependence of the scalar power spectrum. We study the models proposed by Ben-Dayan \& Brustein numerically. For each model, we solve the background eqautions and the MS equations \cite{Mukhanov:1981xt, Sasaki:1983kd, Mukhanov:1985rz} to obtain a primordial power spectrum. This process is then applied to a large sample of models and allows us to numerically study the dependence of the cosmological parameters on the potential parameters with accuracy well below one percent. We found significant differences between the analytical predictions of the commonly used Stewart-Lyth expressions \cite{Stewart:1993bc,Lyth:1998xn} for CMB observables and the numerical results. We discuss these differences and show that when the slow-roll hierarchy is not valid, one cannot rely on the Stewart-Lyth expressions since the time derivatives of the first and second slow-roll parameters ($\epsilon_H,\delta_H$) cannot be neglected. Instead, these models require careful numerical studies to calculate the spectrum. This also means that, in some cases, it is not possible to use Hankel functions as an approximate solution of the Mukhanov-Sasaki (MS) equation. In other cases the Hankel functions can be used, but the order of the Hankel function is different.
We have studied an interesting class of models that can produce a high tensor-to-scalar ratio while conforming to observable values of $n_{s}$ and $n_{run}$. In order to accomplish this, we built a numerical package that calculates the primordial power spectrum from a given slow roll inflationary potential. Conservatively, the errors in the estimated cosmological parameters are bounded from above by $0.1\%$. This package allows us to generate and analyze large number of models. Models that cover the $n_{s}-n_{run}$ region of interest were produced. We set $r_0=0.001$ and probed the dependence of $n_s$ and $n_{run}$ on the potential parameters $\eta_0$ and $\alpha_{0}$, and later studied the behaviour of $n_{s}$ and $n_{run}$ as a function of $r_0$ and $\alpha_{0}$ (setting $\eta_0=0$). We found that the predictions made using the standard Stewart-Lyth analytic expression are not accurate enough and numerical analysis is required in order to study these models accurately. We analysed the origin of this discrepancy in detail, and revealed some approximations that are unjustified for the cases at hand. The small field models discussed in this paper deviate from the standard slow-roll hierarchy and display interesting features associated with the high amplitude of GW. In future work we plan to extend our analysis to models that produce higher values of $r$, compare them to data and determine the best fit candidate for a small field inflationary model.
16
7
1607.03740
We study scale dependence of the cosmic microwave background (CMB) power spectrum in a class of small, single-field models of inflation which lead to a high value of the tensor to scalar ratio. The inflaton potentials that we consider are degree 5 polynomials, for which we precisely calculate the power spectrum, and extract the cosmological parameters: the scalar index $n_s$, the running of the scalar index $n_{\mathrm{run}}$ and the tensor to scalar ratio $r$. We find that for non-vanishing $n_{\mathrm{run}}$ and for $r$ as small as $r=0.001$, the precisely calculated values of $n_s$ and $n_{\mathrm{run}}$ deviate significantly from what the standard analytic treatment predicts. We study in detail, and discuss the probable reasons for such deviations. As such, all previously considered models (of this kind) are based upon inaccurate assumptions. We scan the possible values of potential parameters for which the cosmological parameters are within the allowed range by observations. The 5 parameter class is able to reproduce all of the allowed values of $n_s$ and $n_{\mathrm{run}}$ for values of $r$ that are as high as 0.001. Subsequently this study at once refutes previous such models built using the analytical Stewart-Lyth term, and revives the small field brand, by building models that do yield an appreciable $r$ while conforming to known CMB observables.
false
[ "scalar ratio", "values", "known CMB observables", "building models", "previous such models", "potential parameters", "inaccurate assumptions", "CMB", "n_s$", "the small field brand", "such deviations", "a high value", "the scalar index", "observations", "the allowed values", "small, single-field models", "the possible values", "inflation", "the cosmological parameters", "the precisely calculated values" ]
11.353955
-0.675371
89
12408487
[ "Werner, Klaus", "Rauch, Thomas", "Hoyer, Denny", "Quinet, Pascal" ]
2016ApJ...827L...4W
[ "Detection of Forbidden Line Components of Lithium-like Carbon in Stellar Spectra" ]
3
[ "Institute for Astronomy and Astrophysics, Kepler Center for Astro and Particle Physics, Eberhard Karls University Tübingen, Sand 1, D-72076 Tübingen, Germany", "Institute for Astronomy and Astrophysics, Kepler Center for Astro and Particle Physics, Eberhard Karls University Tübingen, Sand 1, D-72076 Tübingen, Germany", "Institute for Astronomy and Astrophysics, Kepler Center for Astro and Particle Physics, Eberhard Karls University Tübingen, Sand 1, D-72076 Tübingen, Germany", "Physique Atomique et Astrophysique, Université de Mons—UMONS, B-7000 Mons, Belgium ; IPNAS, Université de Liège, Sart Tilman, B-4000 Liège, Belgium" ]
[ "2017ASPC..509..189H", "2018A&A...612A..62H", "2018EL....12363001S" ]
[ "astronomy" ]
4
[ "atomic data", "atomic processes", "stars: atmospheres", "white dwarfs", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1927Natur.120..473B", "1929ApJ....69..173S", "1969A&A.....1...28B", "1969Natur.221..947G", "1974JQSRT..14.1025B", "1974slbp.book.....G", "1976ApJ...204L.119L", "1981tass.book.....C", "1987PhRvA..36.2265B", "1990PhRvA..42.5433C", "1991A&A...244..437W", "1991A&AS...89..581D", "1991JQSRT..45...11A", "1994PhRvE..49.5644G", "1994PhRvE..49.5889G", "1995ApJ...441L..85B", "1998ApJ...496..395B", "2005pps..book.....G", "2014AcSpB.100...86C", "2015A&A...582A..94W", "2016A&A...590A.128R" ]
[ "10.3847/2041-8205/827/1/L4", "10.48550/arXiv.1607.08387" ]
1607
1607.08387_arXiv.txt
\label{intro} Forbidden line components are atomic transitions with $\Delta \ell \not= \pm 1$, where $\ell$ is the angular quantum number. They are associated with the mixing of upper states induced by the plasma electric microfield, leading to transitions that are normally disallowed by the selection rules for electric dipole transitions. This effect should not be confused with forbidden lines associated with magnetic dipole, electric quadrupole or other higher multipole transitions which are well known tools for analysing emission lines from thin astrophysical plasmas, e.g., a multitude of forbidden lines in planetary nebulae \citep{1927Natur.120..473B} and He-like triplets in X-ray spectra of stellar coronae \citep{1969Natur.221..947G}. The forbidden line components investigated here are not restricted to low densities because they do not involve metastable states. In contrast, they appear as absorption lines at high densities when line broadening by the Stark effect is important. \begin{figure*}[bth] \centering\includegraphics[height=0.81\textheight]{fig1a.ps}\hspace{1cm} \includegraphics[height=0.81\textheight]{fig1b.ps} \caption{Observed spectra of DO white dwarfs and PG\,1159 stars (grey lines) with overplotted model spectra (red lines; forbidden components not included). \emph{Left panel:} Forbidden \ion{C}{4} 3p--4f multiplet (indicated by vertical, blue dashed lines) in the blue wing of the allowed 3p--4d line. Object names and spectral types are indicated. \emph{Right panel:} For the same stars, we show the forbidden \ion{C}{4} 3d--4d multiplet in the red wing of the allowed 3d--4f line. \Teff, \logg, and carbon abundance (mass fraction) of the models are indicated. Some other photospheric lines are marked in both panels.\label{fig:all} } \end{figure*} The most prominent example for numerous forbidden components are neutral helium lines in optical spectra of white dwarfs \citep[e.g.,][]{1976ApJ...204L.119L,1995ApJ...441L..85B}, originally detected in B-type stars by \citet{1929ApJ....69..173S} at \ion{He}{1}~$\lambda$4471\,\AA. This $2\,^3P-4\,^3D$ transition is accompanied by the forbidden $\Delta \ell = 2$ component $2\,^3P-4\,^3F$. Detailed descriptions of the physical process associated with the formation of forbidden components in neutral helium can be found, e.g., in \citet{1969A&A.....1...28B,1974JQSRT..14.1025B,1974slbp.book.....G,1991JQSRT..45...11A,2005pps..book.....G}. To our best knowledge, forbidden components of elements heavier than helium have hitherto not been identified in astrophysical plasmas but in laboratory plasmas. \citet{1987PhRvA..36.2265B} report the detection of such components in lithium-like (i.e., one valence electron) \ion{C}{4} and \ion{N}{5}. They can be used for plasma diagnostics because they are strongly density dependent. Another example is lithium itself. \ion{Li}{1} $\lambda$4602\,\AA\ 2p--4d with its forbidden 2p--4p and 2p--4f components is used for diagnostics \citep[e.g.,][]{2014AcSpe.100...86C} in plasmas where helium with its forbidden components is not present for that purpose. The first to investigate the presence of forbidden line components in stellar spectra was, as mentioned, \citet{1929ApJ....69..173S}. He wrote: ``Of the various elements only helium seems to promise any results. Hydrogen shows no new lines outside the Balmer components, which are blended. All other elements are either faint in the stars or not very susceptible to Stark effect.'' At last we can report here on the detection of forbidden line components of a heavier species, namely carbon in stellar spectra. They occur in hot (pre-) white dwarfs, namely, the same two \ion{C}{4} transitions discovered in the plasma experiment by \citet{1987PhRvA..36.2265B}. The detection is favored by the conditions encountered in hot white dwarf atmospheres. Broad \ion{C}{4} lines due to strong Stark effect and, often, highly enriched carbon abundance. Because of the strong density dependence of the line strengths, they can potentially be used as sensitive gravity indicators in stellar spectra. \begin{figure}[bth] \centering\includegraphics[width=0.95\columnwidth]{fig_grotrian.ps} \caption{Grotrian diagram of the lithium-like \ion{C}{4} ion. Solid lines indicate the observed dipole allowed transitions, dashed lines indicate the identified forbidden components.}\label{fig:grotrian} \end{figure}
Appropriate quantum-mechanical calculations for the Stark line broadening of forbidden \ion{C}{4} components are not available. \citet{1994PhRvE..49.5889G} performed such calculations for the interpretation of laboratory spectra with a computer code presented by \citet{1990PhRvA..42.5433C} but, to our best knowledge, they were not published. For the 3p--4f transition, broadening data were published by \citet{1991A&AS...89..581D}, however, they are useless in our context because it was assumed an isolated line. At the moment, we are only able to include the lines by assuming linear Stark effect (as for the allowed components) in the approximation presented by \citet{1991A&A...244..437W} and guessing their strength by arbitrarily upscaling theoretical $f$-values computed by us. As an example we show in Fig.\,\ref{fig:all_det} (black line) the result of this procedure in the case of the DO white dwarf \re. The $f$-values required scaling by factors of 400 and 10\,000 for the 3p--4f and 3d--4d transitions, such that they amount to about $f = 0.004-0.008$. The line positions are matched while the asymmetries are not because, obviously, our assumption for broadening is poor. The extended wings that point away from the allowed line component are not broad enough in the model. An arbitrary increase of the Stark damping constant by a factor of six and a further increase of the $f$-values by a factor of two results in a better fit, but then the steep wings pointing towards the allowed line components are too broad (Fig.\,\ref{fig:all_det}, dashed blue line). The asymmetric line shape of the forbidden components (with their steep wing always towards the adjacent allowed line) is identical to the behaviour of such lines of neutral helium in stellar atmospheres. Under certain circumstances, the asymmetry can be very pronounced as was demonstrated by \citet[][their Fig.\,2]{1998ApJ...496..395B} and was explained by the effect of varying ratios of line widths to the separation of forbidden and allowed components as a function of formation depths of line cores and wings. The detection of the \ion{C}{4} 3d--4d forbidden component at 1171\,\AA\ in the experiment by \citet{1987PhRvA..36.2265B} (observed in second order) was subsequently questioned by \citet{1994PhRvE..49.5644G} who argued that the respective spectral feature is an impurity line, namely the $\lambda$585\,\AA\ \ion{C}{3} 2p\,3s--2p$^2$ line (at 1171\,\AA\ in fourth order). In light of our observations in white dwarfs we conclude that \citet{1987PhRvA..36.2265B} indeed saw the \ion{C}{4} 3d--4d component, at least contributing to the impurity line. The forbidden components of \ion{C}{4} originally identified in laboratory plasmas \citep{1987PhRvA..36.2265B} and now in stellar spectra are $n = 3-4$ transitions, where $n$ is the principal quantum number. The 3d--4d forbidden component of lithium-like \ion{N}{5} was also detected by \citet{1987PhRvA..36.2265B}. It is located at 749\,\AA\ and therefore not accessible in stellar spectra because of extinction by interstellar neutral hydrogen. The respective $n = 3-4$ lines of lithium-like \ion{O}{6} have even shorter wavelengths. \ion{O}{6} has strong and broad lines, comparable to the \ion{C}{4} lines, in the ultraviolet spectra of the PG\,1159 stars presented here. Candidates for forbidden components are $n = 4-5$ transitions, but we could not identify any.
16
7
1607.08387
We report the first identification of forbidden line components from an element heavier than helium in the spectrum of astrophysical plasmas. So far, these components were identified only in laboratory plasmas and not in astrophysical objects. Forbidden components are well known for neutral helium lines in hot stars, particularly in helium-rich post-AGB stars and white dwarfs. We discovered that two hitherto unidentified lines in the ultraviolet spectra of hot hydrogen-deficient (pre-) white dwarfs can be identified as forbidden line components of triply ionized carbon (C IV). The forbidden components (3p-4f and 3d-4d) appear in the blue and red wings of the strong, Stark broadened 3p-4d and 3d-4f lines at 1108 Å and 1169 Å, respectively. They are visible over a wide effective temperature range (60,000-200,000 K) in helium-rich (DO) white dwarfs and PG 1159 stars that have strongly oversolar carbon abundances.
false
[ "white dwarfs", "neutral helium lines", "forbidden line components", "astrophysical objects", "astrophysical plasmas", "hot stars", "unidentified lines", "lines", "C IV", "Forbidden components", "laboratory plasmas", "helium", "helium-rich post-AGB stars", "1108 Å", "1169 Å", "strongly oversolar carbon abundances", "3d-4f", "3p-4d", "helium-rich (DO) white dwarfs", "hot hydrogen-deficient (pre-) white dwarfs" ]
7.843997
9.802689
-1
12623985
[ "Hagala, R.", "Llinares, C.", "Mota, D. F." ]
2017PhRvL.118j1301H
[ "Cosmic Tsunamis in Modified Gravity: Disruption of Screening Mechanisms from Scalar Waves" ]
15
[ "Institute of Theoretical Astrophysics, University of Oslo, PO Box 1029 Blindern, 0315 Oslo, Norway", "Institute for Computational Cosmology, Department of Physics, Durham University, Durham DH1 3LE, United Kingdom", "Institute of Theoretical Astrophysics, University of Oslo, PO Box 1029 Blindern, 0315 Oslo, Norway" ]
[ "2017JCAP...01..031M", "2017PhRvD..95h4029B", "2018IJMPD..2730003M", "2018IJMPD..2748003L", "2018JCAP...06..035I", "2018PhRvD..98f4019O", "2019PhRvD..99l4050K", "2019PhRvL.122i1102L", "2019arXiv190803430B", "2021JCAP...12..043B", "2021PhRvD.103b4009N", "2021PhRvL.127s1101I", "2021RvMP...93a5003B", "2021arXiv210512582S", "2024arXiv240102410C" ]
[ "astronomy", "physics" ]
3
[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Theory" ]
[ "1989RvMP...61....1W", "1998AJ....116.1009R", "2003Natur.425..374B", "2004PhRvL..93q1104K", "2009IJMPD..18.2147B", "2010PhRvL.104w1301H", "2011PhRvL.106y1101S", "2012PhR...513....1C", "2012PhRvL.109x1102K", "2013PhRvL.110p1101L", "2014PhRvD..89h4023L", "2014PhRvD..90b4001B", "2014PhRvD..90l4041L", "2015JCAP...02..034B", "2015PhR...568....1J", "2016A&A...585A..37H", "2016JCAP...11..045B", "2016PhRvD..93d4050L" ]
[ "10.1103/PhysRevLett.118.101301", "10.48550/arXiv.1607.02600" ]
1607
1607.02600_arXiv.txt
16
7
1607.02600
Extending general relativity by adding extra degrees of freedom is a popular approach for explaining the accelerated expansion of the Universe and to build high energy completions of the theory of gravity. The presence of such new degrees of freedom is, however, tightly constrained from several observations and experiments that aim to test general relativity in a wide range of scales. The viability of a given modified theory of gravity, therefore, strongly depends on the existence of a screening mechanism that suppresses the extra degrees of freedom. We perform simulations, and find that waves propagating in the new degrees of freedom can significantly impact the efficiency of some screening mechanisms, thereby threatening the viability of these modified gravity theories. Specifically, we show that the waves produced in the symmetron model can increase the amplitude of the fifth force and the parametrized post Newtonian parameters by several orders of magnitude.
false
[ "extra degrees", "gravity", "freedom", "high energy completions", "such new degrees", "several orders", "several observations", "these modified gravity theories", "general relativity", "magnitude", "scales", "a screening mechanism", "some screening mechanisms", "a given modified theory", "the extra degrees", "Universe", "experiments", "the parametrized post Newtonian parameters", "Newtonian", "the new degrees" ]
9.962138
1.119478
-1
12473385
[ "Xie, Fu-Guo", "Zdziarski, Andrzej A.", "Ma, Renyi", "Yang, Qi-Xiang" ]
2016MNRAS.463.2287X
[ "A luminous hot accretion flow in the low-luminosity active galactic nucleus NGC 7213" ]
16
[ "Key Laboratory for Research in Galaxies and Cosmology, Shanghai Astronomical Observatory, Chinese Academy of Sciences, 80 Nandan Road, Shanghai 200030, China", "Nicolaus Copernicus Astronomical Center, Polish Academy of Sciences, Bartycka 18, PL-00-716 Warszawa, Poland", "Department of Astronomy and Institute for Theoretical Physics and Astrophysics, Xiamen University, Xiamen, Fujian 361005, China", "Key Laboratory for Research in Galaxies and Cosmology, Shanghai Astronomical Observatory, Chinese Academy of Sciences, 80 Nandan Road, Shanghai 200030, China" ]
[ "2017ApJ...836..104X", "2017MNRAS.471.2848L", "2018ApJ...860..134Q", "2018MNRAS.475.1190Y", "2019MNRAS.490.4606B", "2019arXiv191200164A", "2020ApJ...890L..29A", "2020ApJ...891...31X", "2020MNRAS.492.2335L", "2020RAA....20..122W", "2021MNRAS.506.4188L", "2022ApJ...926..209S", "2022ApJ...931..167N", "2023MNRAS.519.3903L", "2023MNRAS.521.2215S", "2024ApJ...965..128Y" ]
[ "astronomy" ]
6
[ "accretion", "accretion discs", "galaxies: active", "galaxies: individual: NGC 7213", "galaxies: Seyfert", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1973A&A....24..337S", "1979ApJ...232...34B", "1984ApJ...285..475H", "1989MNRAS.240..833T", "1994ApJ...428L..13N", "1994MNRAS.268..405N", "1995ApJ...452..710N", "1996ApJ...471..106F", "1997ApJ...489..865E", "2000ApJ...542..235K", "2000rdgr.conf....5G", "2001AIPC..599..991T", "2001MNRAS.324..119Y", "2001MNRAS.328..393A", "2002ApJ...564..120H", "2003A&A...407L..21B", "2003ApJ...594L..99Y", "2003MNRAS.342..355Z", "2003MNRAS.343L..59H", "2003MNRAS.345.1057M", "2004A&A...414..895F", "2004MNRAS.354..953Y", "2005ApJ...620..905Y", "2005ApJ...629..408Y", "2005MNRAS.356..727S", "2005MNRAS.356..734B", "2006ApJ...636...56G", "2006MNRAS.368..903P", "2006PASJ...58..193M", "2007ApJ...659..541Y", "2008ARA&A..46..475H", "2008MNRAS.389L..52B", "2009ApJ...699..626H", "2009MNRAS.394.1640S", "2009MNRAS.399..349G", "2009MNRAS.399.1597S", "2010A&A...515A..23H", "2010ApJ...712..639O", "2010MNRAS.408..551L", "2011A&A...536A..36A", "2011ApJ...729...19K", "2011MNRAS.411..402B", "2011MNRAS.413.2259S", "2011MNRAS.414..677C", "2012ApJ...761..129Y", "2012ApJ...761..130Y", "2012MNRAS.423..663Z", "2012MNRAS.423L..87M", "2012MNRAS.424.1327E", "2012MNRAS.427.1580X", "2012agn..book.....B", "2013ApJ...774L..25K", "2013ApJ...777..102Q", "2013MNRAS.428.2500C", "2013MNRAS.429.3439E", "2014A&A...562A.142M", "2014ARA&A..52..529Y", "2014ApJ...781L...2A", "2014ApJ...788...52C", "2014MNRAS.437..316K", "2014MNRAS.438.3322S", "2014MNRAS.438.3434R", "2014MNRAS.440.2238Z", "2014MNRAS.442L.110X", "2015A&A...584A..20W", "2015ApJ...801...47Y", "2015ApJ...804..101Y", "2015MNRAS.447.1692Y", "2015MNRAS.448.1099Q", "2015MNRAS.452.3266U", "2016MNRAS.456.4377X", "2016MNRAS.459.1543W", "2016MNRAS.459.3963C", "2017MNRAS.466..705S" ]
[ "10.1093/mnras/stw2132", "10.48550/arXiv.1607.00928" ]
1607
1607.00928_arXiv.txt
\label{sec:intro} There is a general consensus that Seyferts, low-ionization nuclear emission-line regions (LINERs), radio galaxies and some transition galaxies contain supermassive black holes (BHs) accreting at moderately low accretion rates \citep{Ho08}. We call them here low-luminosity active galactic nuclei (LLAGNs). They have the bolometric luminosity of $L_{\rm bol}\la (0.01$--$0.02)\ \ledd$, where the Eddington luminosity (for the H fraction of 0.7) is $\ledd\simeq 1.47\times 10^{46} \ergs\ (\mbh/ 10^8 \msun)$ and $\mbh$ is the BH mass. LLAGNs have properties distinctly different from bright AGNs, see \citet{Ho08} for a review. In particular, the spectra of LLAGNs generally do not have big blue bumps, but instead peak in the mid-infrared \citep{Ho09}. Also, the optical--X-ray spectral index, $\alpha_{\rm OX}$ \citep{Sobolewska09,Sobolewska11}, the radio loudness \citep*{Ho02, Greene06} and the X-ray photon index, $\Gamma$ (\citealt{Gu09,Emm2012}, hereafter E12; \citealt{Yang2015, Connolly2016}), are all observed to anti-correlate with the Eddington ratio (reported either for the bolometric or X-ray band) in LLAGNs, while they correlate positively in bright AGNs. These differences show that LLAGNs are not simply bright AGNs scaled down in the accretion rate. The leading theoretical picture of LLAGNs is the accretion--jet scenario. This model utilizes three components \citep*{Esin1997, YCN2005, Ho08, Yuan2014}, namely a cold geometrically-thin Shakura-Sunyaev disc (SSD; \citealt{SS1973}) truncated at certain transition radius, $R_{\rm tr}$, a hot accretion flow within this radius, and a jet. The hot component can either be an advection-dominated accretion flow (ADAF; \citealt{Narayan1994}) when the accretion rate, $\mdot$, is low or a luminous hot accretion flow (LHAF; \citealt{Yuan2001, Yuan2003}) when $\mdot$ is intermediate. The accretion-jet model has been successfully applied to both LLAGNs and the hard state of accreting BH binaries (BHBs), see \citet{Yuan2014} for a recent review. We study here NGC~7213, a nearby face-on Seyfert 1/LINER with a number of interesting properties. The source redshift is $z_{\rm r}=0.005839$, and its BH mass is estimated as $\mbh=9.6^{+6.1}_{-4.1}\times 10^7\ \msun$ \citep*{Blank2005}. Using $H_0=67.8$ km s$^{-1}$ Mpc$^{-1}$ \citep{planck} and the redshift corrected to the reference frame defined by the microwave background\footnote{\url{http://ned.ipac.caltech.edu}} of 0.005145, the luminosity distance, $d_{\rm L}$, is 22.8 Mpc. The bolometric luminosity has been estimated as $L_{\rm bol} \simeq 9 \times 10^{42}$ and $\simeq 1.8 \times 10^{43}\ergs$ by \citet{SSNF2014} and E12, respectively, though the latter estimate may contain an infrared (IR) contribution from the local starburst. At the best-fitting mass, those values correspond to $\simeq 0.6\times 10^{-3}$ and $\simeq 1.3 \times 10^{-3}\,L_{\rm Edd}$, respectively. The monochromatic radio (defined as $\lr=\nu L_\nu$ at, e.g., 5 or 8.5 GHz) luminosity, the 2--10 keV X-ray luminosity and the BH mass in AGNs and BHBs were found to be relatively tightly correlated \citep*{Merloni2003, Falcke2004}, forming the so-called Fundamental Plane (hereafter FP). The FP can be expressed as $\lr\propto \lx^p\ \mbh^q$ with the indices of $p=0.6\pm 0.1,\, q=0.78\pm 0.11$ \citep{Merloni2003}. The case of BHBs, where there is almost no dependence on the BH mass, has been reviewed by \citet{Corbel2013}. Physically, the FP provides evidence for a strong connection between the hot X-ray emitting source, and the radio source, usually a jet. In the accretion-jet model, the radio is from the jet, and unless the source is sufficiently faint the X-rays will emit from the hot accretion flow. The FP is interpreted as the consequence of a tight coupling between accretion and ejection, which is most evident in the mass flow ratio $\eta_{\rm jet}$ (cf. Sec. \ref{sec:rxmodel} for definition), see \citet{YC2005} and \citet{Xie2016} for details. \citet{Bell2011}, herafter B11, carried out intense monitoring of NGC~7213 in radio and X-rays over three years. They found a significant correlation, but with a clear deviation in its slope with respect to the FP. Similar deviations are also observed in a few other AGNs (NGC~4051 and NGC~4395; \citealt{King2011, King2013}) and BHBs, where they have been named outliers \citep{Coriat2011, Corbel2013, Xie2016}. The latter seem to follow a hybrid correlation, most evident in H1743--322 \citep{Coriat2011}, which is the outlier with the largest dynamic range in $\lx$. The radio/X-ray correlation of H1743--322 exhibits three branches \citep{Coriat2011}: it is steep, $p\approx 1.3$, at large $\lx$, flat, $p\sim 0$, at moderate $\lx$, and has the original FP slope, $p\approx 0.6$, at low $\lx$. \begin{figure*} \centering \includegraphics[width=13.cm]{ngc7213-sed_w_fermi.eps} \caption{The broad-band SED of NGC~7213. The red dotted, blue dot--dashed and green dashed curves show the emission from the truncated SSD, the LHAF and the jet, respectively. The black solid curve gives the total model. The radio to UV data are from E12, while the X-ray (0.6--150 keV) data, shown by the thick red line, are from \citet{Lobban2010}. The \fermi/LAT data are taken from \citet{Woja2015}. } \label{sed} \end{figure*} \begin{figure} \centering \includegraphics[width=8.cm]{ngc7213-rxcorrnew.eps} \caption{The radio/X-ray correlation of NGC~7213, based on B11. The grey filled circles are for 4.8 GHz versus the 2--10 keV luminosity during MJD $\simeq$53800--54900 (2006 July--2009 September), which have been shifted by the 35-d lag of radio with respect to X-rays (as in fig. 7 of B11). The green filled pentagon and red filled squares show archival data (MJD 47527--52057) collected by B11, for which the average interval between the radio (at 8.4 GHz) and X-ray measurements was $\simeq$1.5 yr. As B11, the errorbars of these archival data are estimated as the fraction rms scatter of the 3 yr ATCA/\xte\ data. The two black long-dashed curves represent the fitting of the grey circles and the red squares, respectively. Our model results are shown by the magenta crosses connected by solid line and the blue crosses connected by dashed line, see Section \ref{sec:rxmodel} for details. They approximately follow relation of $\lr\propto \lx^{0.15}$ and $\lr\propto \lx^{1.3--1.4}$, respectively. For comparison, the FP fit for $\mbh=10^8\, \msun$ \citep{Merloni2003} is shown by the grey dashed line.} \label{rx_corr} \end{figure} Here, we point out that the radio/X-ray correlation in NGC~7213 resembles that of H1743--322. \citet{Xie2016} recently investigated both the FP-type and the hybrid radio/X-ray correlations in BHBs under the accretion-jet model. In their work, systems with the FP relationship were postulated to stay in the ADAF accretion regime, while systems with hybrid radio/X-ray correlation would enter different hot accretion modes. In the latter case, they postulated that an LHAF is responsible for the flat branch (with $p\sim0$), and a two-phase accretion flow (also called type II LHAF) is responsible for the steep ($p\sim 1.3$) branch. We develop these ideas here for NGC~7213, arguing that its central engine also enters the LHAF regime. We note that alternative interpretations exist for the steep branch (at high $\lx$). \citet{MM2014}, \citet*{Cao2014} and \citet{Qiao2015} interpreted it in terms of the disc-corona model, while \citet*{Yu2015} proposed that it is due to a radiatively-cooled magnetized hot accretion flow. However, none of those models can explain the flat transition branch, which apparently represents a connection between the FP-type correlation and the steep one. This paper is organized as follows. Section \ref{sec:obs} describes main observational properties of NGC~7213. Section \ref{sec:model} describes the LHAF-jet model, which is then applied to NGC~7213 in Section \ref{sec:application}. This allows us to understand its broad-band radio-to-$\gamma$-ray spectrum, time lags of the radio emission with respect to the X-rays, and constraints on the size of the radio source. In Section \ref{compare}, we compare NGC 7213 to other sources. Section \ref{sec:summary} gives our conclusions.
\label{sec:summary} Our main results are as follows. We compile an exhaustive set of the observational properties of NGC~7213. These include its average broad-band spectrum, the relationships between its X-ray luminosity with the radio luminosity and the X-ray spectral index covering a broad range of $\lx$, and the radio properties. We successfully model these properties through the accretion-jet model, which includes three components, i.e. a truncated cold disc, an inner hot flow and a jet. The inner hot flow changes its mode from the LHAF to two-phase accretion, which explains the changes of the slopes in the radio/X-ray and index/flux correlations. The critical 2--10 keV luminosity at which the change of accretion mode occurs is $L_{\rm X, LHAF, crit} \approx 1.5 \times10^{42} \ergs \approx 1.0\times 10^{-4} \ledd$. The model parameters of the composite SED with $\lx\lesssim L_{\rm X, LHAF, crit}$ are listed in Table 1. The most important parameter determined by us is the viscosity parameter, which is found to be relatively small. More explicitly, we have $\alpha\approx 0.011\ (\mbh/10^8\ \msun)^{-1}\ (d_{\rm L}/22.8\ {\rm Mpc})^2$, where the uncertainties in $\mbh$ and/or $d_{\rm L}$ are taken into account. This low $\alpha$ value allows the accretion flow to be in the LHAF regime at a relatively low bolometric luminosity. This determination is in agreement with \citet{Xie2016}, who proposed that systems with hybrid radio/X-ray correlation have small $\alpha$ values. Based on the observed time lags, the radio flux and the radio size constraint, we find the jet Lorentz factor to be small, $\Gamma\la 4$. The location of the radio source at the observed $\nu=8.4$ GHz is $\sim 10^4 R_{\rm g}$.
16
7
1607.00928
The active galactic nucleus (AGN) NGC 7213 shows a complex correlation between the monochromatic radio luminosity L<SUB>R</SUB> and the 2-10 keV X-ray luminosity L<SUB>X</SUB>, I.e. the correlation is unusually weak with p ∼ 0 (in the form L_R∝ L_X^p) when L<SUB>X</SUB> is below a critical luminosity, and steep with p &gt; 1 when L<SUB>X</SUB> is above that luminosity. Such a hybrid correlation in individual AGNs is unexpected as it deviates from the Fundamental Plane of AGN activity. Interestingly, a similar correlation pattern is observed in the black hole X-ray binary H1743-322, where it has been modelled by switching between different modes of accretion. We propose that the flat L<SUB>R</SUB>-L<SUB>X</SUB> correlation of NGC 7213 is due to the presence of a luminous hot accretion flow, an accretion model whose radiative efficiency is sensitive to the accretion rate. Given the low luminosity of the source, L<SUB>X</SUB> ∼ 10<SUP>-4</SUP> of the Eddington luminosity, the viscosity parameter is determined to be small, α ≈ 0.01. We also modelled the broad-band spectrum from radio to γ-rays, the time lag between the radio and X-ray light curves, and the implied size and the Lorentz factor of the radio jet. We predict that NGC 7213 will enter into a two-phase accretion regime when L<SUB>X</SUB> ≳ 1.5 × 10<SUP>42</SUP> erg s<SUP>- 1</SUP>. When this happens, we predict a softening of the X-ray spectrum with the increasing flux and a steep radio/X-ray correlation.
false
[ "γ-rays", "accretion", "L", "X</SUB", "ray", "-", "radio", "the X-ray spectrum", "a steep radio/X-ray correlation", "the radio and X-ray light curves", "AGN activity", "a luminous hot accretion flow", "the monochromatic radio luminosity L<SUB>R</SUB", "the accretion rate", "an accretion model", "different modes", "the 2-10 keV X-ray luminosity L<SUB>X</SUB", "AGN", "the radio jet", "the flat L<SUB>R</SUB>-L<SUB>X</SUB> correlation" ]
15.479712
7.255192
119
12436094
[ "Chamel, N." ]
2016arXiv160705943C
[ "Constraint on the internal structure of a neutron star from Vela pulsar glitches" ]
2
[ "-" ]
[ "2018arXiv181001743P", "2021MNRAS.500.5336W" ]
[ "astronomy" ]
3
[ "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "2005AJ....129.1993M", "2007NuPhA.789..403D", "2009ASSL..357.....B", "2011MNRAS.414.1679E" ]
[ "10.48550/arXiv.1607.05943" ]
1607
1607.05943_arXiv.txt
Pulsars, the compact remnants of gravitational core-collapse supernova explosions, are spinning very rapidly with extremely stable periods $P$ ranging from about 1.4~milliseconds to a few seconds~\cite{beck09}. The variations $\dot{P}\equiv dP/dt$ of the rotation period of some pulsars over time $t$ do not exceed $10^{-21}$, as compared to $10^{-18}$ for the most accurate atomic clocks~\cite{hinkley2013}. Still, some pulsars have been found to suddenly spin up. The ensuing ``glitches'' in their rotational frequency $\Omega$, ranging from $\Delta\Omega\slash \Omega\sim 10^{-9}$ to $\sim 10^{-5}$, are generally followed by a long relaxation lasting from days to years, and sometimes accompanied by an abrupt change of the spin-down rate from $\vert\Delta \dot\Omega\slash\dot\Omega\vert\sim 10^{-6}$ up to $\sim 10^{-2}$. At the time of this writing, 472 glitches have been detected in 165 pulsars~\cite{jod12}. One of the most emblematic glitching pulsars is Vela (PSR B0833$-$45). The Vela pulsar was discovered in October 1968 by astronomers from Sydney University~\cite{large68}. Its very short period of about 89 milliseconds provided very strong evidence for the identification of pulsars as rotating neutron stars rather than rotating or vibrating white dwarfs. Moreover, the fact that this pulsar lies in a supernova remnant confirmed the scenario of gravitational core collapse of massive stars proposed much earlier by Baade and Zwicky~\cite{bz34}. Between 24 February and 3 March 1969, the Vela pulsar was found to spin more rapidly than before~\cite{rad69,rei69}: the rotational frequency had increased by $\Delta \Omega/\Omega \simeq 2\times 10^{-6}$. The increase in the spin-down rate was even larger $\Delta \dot\Omega/\dot\Omega\simeq 7\times 10^{-3}$. Soon afterwards, Malvin Ruderman proposed that the glitch\footnote{According to George Greenstein, the term ``glitch'' was coined by Ruderman~\cite{green83}.} was the manifestation of crustquakes~\cite{rud69}. The idea was the following. As the star spins down, the presence of a solid crust prevents any readjustment of the stellar shape. Stresses build up to the point at which the solid crust will break down. The resulting quake entails a sudden decrease in the moment of inertia, and an increase in the spin frequency by conservation of angular momentum. As David Pines pointed out in 1999~\cite{pines99}, the Vela quake would be a cataclysmic event since this would correspond to an Earthquake of magnitude 17 in which the entire surface is shifted by about 15 meters! On the other hand, crustquakes should be very rare: this model predicted that no such event should be observed again in Vela in a human lifetime~\cite{smo70,baym71}. The detection of the first glitch a few months after the discovery of the Vela pulsar thus appeared as a very unlikely coincidence, as remarked by Ruderman himself~\cite{rud69}: ``The proposed model does not, however, account for the remarkable fact that such an uncommon event should just happen to occur during the short time, less than a year, during which this pulsar has been observed.'' Freeman Dyson speculated that neutron-star quakes could be more frequent if one assumes that the crustal stress arises from the accumulation of volcanic ashes at the stellar surface~\cite{dyson69}. In the fall of 1971, a second glitch occurred in Vela thus definitively ruling out Ruderman's crustquake theory. Other scenarios were proposed like corequakes~\cite{pines72}, planetary perturbation~\cite{michel70} or magnetospheric instabilities~\cite{scargle71}, but none of them were really convincing (see, e.g. Ref.~\cite{pines74}). For instance, if Vela glitches arose from corequakes they should have been seen in X-rays~\cite{rud76}. The eccentric orbital motion of planets around Vela predicted a strict regularity of glitches that were later found to be inconsistent with observations. Finally, the absence of radiative and pulse profil changes during glitches ruled out a magnetospheric origin. Since no such phenomena had ever been observed in other celestial bodies, glitches had to do with specific properties of neutron stars. Cameron and Greenstein proposed that the Vela glitch was due to the onset of fluid instabilities, the core rotating faster than the crust assuming ``viscous effects to be unimportant''~\cite{cam69}. After one year, the spin-down rate of Vela relaxed to the value it had before the glitch. This very long relaxation time was interpreted as a strong evidence for neutron-star superfluidity~\cite{baym69}. Neutron-star superfluidity had been actually predicted long before the discovery of pulsars by Migdal in 1959~\cite{mig59}, and had been first studied by Ginzburg and Kirzhnits in 1964~\cite{gk64}. In 1972, Packard suggested that glitches are related to the metastability of this superfluid~\cite{pac72}. In 1975, Anderson and Itoh advanced the seminal idea that glitches are triggered by the sudden unpinning of superfluid vortices in neutron-star crust~\cite{and75}. A rotating superfluid is threaded by a regular array of quantized vortex lines, each carrying a quantum $\hbar$ of angular momentum. Such vortices have been observed in various kinds of superfluids in the laboratory and are thus expected to be present in the interior of rotating superfluid neutron stars. Quantized vortices tend to arrange themselves on a regular triangular array. The surface density of vortices in a rotating fermionic superfluid is given by $\displaystyle n_v=2m\Omega/(\pi\hbar)$, where $m$ is the mass of the fermions, $\Omega$ the angular frequency and $\hbar$ the Planck-Dirac constant. For a neutron superfluid, we obtain $\displaystyle n_v ({\rm km}^{-2}) \sim 10^{14} /P({\rm s})$. The Vela pulsar is thus expected to contain about $10^{17}$ vortices. The neutron superfluid in a neutron star is supposed to be weakly coupled to the rest of the star by mutual friction forces and to thus follow its spin-down via the motion of vortices away from the rotation axis unless vortices are pinned to the crust. In this case, the superfluid could rotate more rapidly than the crust. The lag between the superfluid and the crust induces a Magnus force acting on the vortices thereby producing a crustal stress. At some point, the vortices will suddenly unpin, the superfluid will spin down and, by the conservation of angular momentum, the crust will spin up leading to a glitch. This scenario found support from laboratory experiments with superfluid helium~\cite{tsa80}. Due to (non-dissipative) mutual neutron-proton entrainment effects, neutron superfluid vortices in the core of a neutron star carry a fractional magnetic quantum flux~\cite{sed80}. As a consequence, the core neutron superfluid is strongly coupled to the crust due to electrons scattering off the magnetic field of the vortex lines~\cite{alp84}. For this reason, only the neutron superfluid permeating the inner crust of a neutron star have been generally thought to be responsible for Vela glitches. Alpar, Pines and collaborators extended this vortex-mediated glitch scenario to explain the postglitch relaxation by the creeping of vortices~\cite{alp85,alp93}. Ruderman argued that if the pinning is strong enough the solid crust will crack before vortices get unpinned~\cite{rud91} (see also Chapter 14 of Ref.~\cite{beck09}). The subsequent motion of crustal plates could naturally explain the observed increase of the spin-down rate. Carter and collaborators later showed that the lack of centrifugal buoyancy and stratification are sources of crustal stress even in the absence of vortex pinning~\cite{car00,cc06}. Whether or not glitches are triggered by crust quakes remains uncertain (see, e.g. Refs.~\cite{lrr,hask15} and references therein for a review of more recent developments). Regardless of the actual glitch triggering mechanism, Vela pulsar timing data have been used to constrain the crustal moment of inertia of a neutron star~\cite{alp93,dat93,lnk99}. The fact that the inferred crustal moment of inertia is $\sim 2-3\%$, as expected for a neutron star with a typical mass $M\sim 1.4 M_\odot$ ($M_\odot$ being the mass of the Sun), provided additional support to the vortex-mediated glitch scenario. On the other hand, this scenario has been recently challenged~\cite{and12,cha13,hoo15,new15} by the realization that despite the absence of viscous drag the crust can still resist the flow of the neutron superfluid due to Bragg scattering~\cite{cha04,cha05,cch05,cha12}. In this paper, the extent to which the crustal superfluid can explain Vela pulsar glitches given the current lack of knowledge of nuclear physics is more closely examined considering a set of different unified dense-matter equations of state, treating consistently all regions of a neutron star, and taking into account nuclear-physics uncertainties.
The giant frequency glitches observed in the Vela pulsar provide strong evidence for the presence of superfluids in the interior of neutron stars. Predicted before the discovery of pulsars, neutron-star superfluidity has recently found additional support from the rapid cooling of the neutron star in Cassiopeia A~\cite{pag11,sht11}, from observations of the initial cooling in persistent soft X-ray transients~\cite{shternin07,brown09}, and from measurements of pulsar braking indices~\cite{alp06,ho12} (for a recent review of neutron-star superfluidity, see, e.g. Ref.~\cite{page14}). According to the standard theory~\cite{alp85}, Vela pulsar glitches are triggered by the sudden unpinning of neutron superfluid vortices in the crust of the star. Assuming that this scenario is correct, a constraint on the crustal moment of inertia of Vela can be inferred from timing data~\cite{dat93,alp93,lnk99}. Although giant glitches have been observed in other pulsars, observations of the Vela pulsar yield the most stringent constraint: $I_{\rm crust}/I\geq 1.6\%$. This constraint has been recently revised after realizing that the crustal superfluid is strongly entrained by the rest of the star due to Bragg scattering~\cite{cha13}: $I_{\rm crust}/I\geq 9.3\%$. Considering current experimental and theoretical constraints on the dense-matter equation of state, we find that the crustal superfluid does not carry enough angular momentum to explain Vela pulsar glitches. Even if crustal entrainment is ignored, the standard vortex-mediated glitch scenario has been challenged by the observation in 2010 of a huge glitch in PSR~2334$+$6 from which the following constraint was inferred: $I_{\rm crust}/I\geq 9.4\%$~\cite{alp11}. These observations suggest that the neutron superfluid in the core of a neutron star contributes to glitches~\cite{gug14}.
16
7
1607.05943
Pulsars are spinning extremely rapidly with periods as short as about $1.4$ milliseconds and delays of a few milliseconds per year at most, thus providing the most accurate clocks in the Universe. Nevertheless, sudden spin ups have been detected in some pulsars like the emblematic Vela pulsar. These abrupt changes in the pulsar's rotation period have long been thought to be the manifestation of a neutron superfluid permeating the inner crust of neutron stars. However, the neutron superfluid has been recently found to be so strongly coupled to the crust that it does not carry enough angular momentum to explain the Vela data. We explore the extent to which pulsar-timing observations can be reconciled with the standard glitch theory considering the lack of knowledge of the dense-matter equation of state.
false
[ "state", "enough angular momentum", "Pulsars", "neutron stars", "Vela", "the emblematic Vela pulsar", "knowledge", "the Vela data", "periods", "year", "pulsar-timing observations", "delays", "sudden spin ups", "the standard glitch theory", "a neutron superfluid", "the neutron superfluid", "a few milliseconds", "the dense-matter equation", "the inner crust", "some pulsars" ]
5.179785
3.237667
69
12585788
[ "Stürmer, Julian", "Seifahrt, Andreas", "Schwab, Christian", "Bean, Jacob L." ]
2017JATIS...3b5003S
[ "Rubidium-traced white-light etalon calibrator for radial velocity measurements at the cm s<SUP>-1</SUP> level" ]
16
[ "University of Chicago, Department of Astronomy and Astrophysics, Chicago, Illinois, United States", "University of Chicago, Department of Astronomy and Astrophysics, Chicago, Illinois, United States", "Australian Astronomical Observatory, Sydney, Australia", "University of Chicago, Department of Astronomy and Astrophysics, Chicago, Illinois, United States" ]
[ "2017JAI.....650001B", "2017PASP..129c4002B", "2018MNRAS.479..768P", "2018SPIE10702E..6DS", "2018SPIE10702E..6TB", "2018SPIE10706E..29H", "2019PASP..131l4503C", "2020SPIE11203E..1NR", "2020SPIE11447E..1FS", "2020arXiv200313770J", "2021arXiv210308456T", "2022A&A...664A.191S", "2022AJ....163..168W", "2022SPIE12184E..1GS", "2022arXiv220405713S", "2024ApOpt..63D..14B" ]
[ "astronomy" ]
5
[ "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1987Metro..23..145B", "2003Msngr.114...20M", "2005OptLE..43..291A", "2007MNRAS.380..839M", "2008Natur.452..610L", "2009ApPhB..96..251S", "2009SPIE.7440E..0MW", "2010MNRAS.405L..16W", "2010RScI...81f3105Q", "2010SPIE.7735E..4XW", "2010Sci...330..653H", "2011A&A...528A.112L", "2011A&A...534A..58P", "2011ApJ...726...73H", "2011SPIE.8151E..1FW", "2012OExpr..20.6631Y", "2012OExpr..2013711P", "2012SPIE.8446E..8EW", "2014A&A...569A..77R", "2014PASP..126..445H", "2014SPIE.9147E..15A", "2014SPIE.9147E..1FQ", "2014SPIE.9147E..3LM", "2014SPIE.9147E..7MG", "2015A&A...581A.117B", "2015ApJ...799...87B", "2015PASP..127..880S", "2016ApJ...816...95G", "2016ExA....42..285F", "2016PASP..128f6001F", "2016SPIE.9908E..18S", "2016SPIE.9908E..70G", "2016SPIE.9912E..29S" ]
[ "10.1117/1.JATIS.3.2.025003", "10.48550/arXiv.1607.05172" ]
1607
1607.05172_arXiv.txt
\label{sec:intro} The radial velocity method has been one of the most important observational techniques in the history of the field of exoplanet science, and it will continue to be critical for making many of the most significant exoplanet discoveries anticipated over the next two decades. One reason for this is that radial velocity surveys can efficiently reveal the census, masses, and orbits of objects in the planetary systems around nearby stars \cite{howard10, mayor11}. These data are valuable constraints on theories of planet formation and evolution. Furthermore, knowing the orbits of the exoplanets around the nearest stars will also be useful for guiding future direct imaging efforts to obtain their spectra \cite{brown15}. Another reason for the ongoing importance of the radial velocity method is that radial velocity measurements are a way to confirm and measure the masses of planet candidates identified by transit searches and to search for additional non-transiting planets in the same systems \cite{gettel16}. Transiting planets are the only planets for which we can determine masses and radii, and examine their atmospheres. Knowing a planet's mass and radius together give us constraints on its bulk composition, and spectroscopic studies can reveal further details about the composition and physical conditions of its atmosphere. Together, bulk composition and atmospheric properties are a powerful diagnostic of a planet's origins and evolution. One of the main goals of the field of exoplanet science is to ascertain the frequency, and the physical and orbital characteristics of low-mass planets. The identification and characterization of potentially habitable worlds is an aspect of this quest that is especially exciting. The techniques for Doppler spectroscopy have currently progressed to the point that precisions of 1 -- 2\,m\,s$^{-1}$ are routinely obtained on bright stars \cite{howard11, lovis11}, and precisions of 60 -- 80\,cm\,s$^{-1}$ have been obtained in a few select cases \cite{pepe11}. This level of precision enables the detection of Earth-mass planets with periods up to a few tens of days around solar-type stars. However, Earth imparts a radial velocity signal with a semi-amplitude of only 9\,cm\,s$^{-1}$ on the Sun when viewed edge on. Therefore, the detection of terrestrial planets in the habitable zones of Sun-like stars is still out of reach with current instruments, and the precision achievable with the radial velocity technique must be substantially improved to pursue this compelling goal. One of the key instrumental aspects that must be improved in radial velocity measurements to reach the level of precision required for the detection of Earth analogues is the spectrograph wavelength calibration \cite{fischer16}. Following the stunning success of the HARPS instrument \cite{mayor03}, the next generation of radial velocity instruments are being designed to feature high thermo-mechanical stability. However, even at the highest levels of instrumental stability, long-term radial velocity drifts exceeding the signal for an Earth-like planet are to be expected. Therefore, the next generation of radial velocity spectrographs require a wavelength calibrator that provides both the local information density and the long-term stability to measure and track instrumental drifts with a precision of \SI{10}{\cmps} or better over time scales of minutes to years. In addition, detailed instrument characterization, e.g. to determine the stitching error of the \replaced{detector}{CCD} \cite{wilken10}, is required for sub-\si{\mps} RV measurements. All these requirements lead to the necessity for the calibration source to provide a dense grid of evenly distributed lines of uniform intensity in order to maximize the extraction of the Doppler information content over a broad wavelength range. As of today the technology most closely fulfilling the requirements for next generation wavelength calibration is the laser frequency comb (LFC) generated by femtosecond pulsed lasers \cite{murphy07,li08,quinlan10, ycas12, phillips12}. The self-referencing of the laser-comb lines produces exceptional stability and frequency precision (better than $10^{-14}$), locked to an atomic standard for very high accuracy. A drawback of LFCs is that the comb lines have a frequency spacing of only a few 100\,MHz. Astronomical spectrographs used for radial velocity measurements usually have a resolution of 10--30\,GHz, which is insufficient to resolve the narrowly spaced comb structure. To increase the comb line spacing, a common solution is to filter out most of the lines with external Fabry-P\'erot (FP) cavities \cite{steinmetz09,quinlan10}, creating a so-called ``astro-comb''. Over the last decade, LFCs for astronomical applications have seen a tremendous development, culminating in the recent introduction of a commercial product offered by Menlo Systems (Germany), covering 450--1400\,nm at 25\,GHz line spacing and $<$3dB intensity variations over the entire bandpass. Precursors of this system have been developed for HARPS-S and ESPRESSO \cite{locurto12m}. However, the cost for this system is often prohibitive, and despite its advantages recent RV spectrograph projects, e.g. CARMENES/CAHA3.5m \cite{carmenes}, SPIROU/CFHT \cite{spirou}, or \replaced{KPF}{Shrek}/Keck \added[]{\cite{gibson2016}}, either do not use an LFC or include it only as a future upgrade, owing to the cost impact on the project. Therefore, a cheaper alternative for wavelength calibration of next-generation RV instruments is desirable. A simple and elegant alternative approach to generating a comb reference spectrum is to use the cavity resonance lines of a Fabry-P\'erot (FP) etalon illuminated with white light \cite{wildi09,wildi10,wildi11,wildi12,halverson12,halverson14}. An ``ideal'' FP etalon illuminated with a broadband light source produces emission lines that are equidistant in frequency space. The position, line width, spacing, and amplitude of these synthetic lines can be easily tailored to match the spectrograph requirements by tuning the finesse and free spectral range of the etalon, which are functions of the separation and reflectivity of the cavity mirrors. The position of the peaks in an etalon spectrum is only a function of cavity length $l$ (essentially mirror spacing $d$, illumination angle $\Theta$, and refractive index $n$) and interference order $m$. In contrast to the LFC, the absolute frequency of the peaks of the etalon comb are not known with great accuracy, owing to the degeneracy between $l$ and $m$. However, what is important is the stability of the line positions, and this can be achieved by keeping the effective cavity length of the etalon constant. It is the lack of the extreme accuracy with which the frequency of each and every comb line is known, that makes the FP etalon inferior, but at the same time so much cheaper than a LFC. Fortunately, in the context of radial velocity measurements, seeking relative RV measurements with extreme precision, extreme stability is required, not extreme global accuracy, and the FP etalon can fulfill this requirement. Practically, however, the stability of a passively stabilized etalon can not be achieved over the desired time scales of exoplanet signals of weeks to years. Measuring and tracking of the etalon line positions with high precision becomes mandatory and we demonstrate the end-to-end implementation and performance of this method in this paper. The FP etalon-based wavelength calibration system described here was developed for the MAROON-X instrument, which is a new fiber-fed, red-optical, high-precision radial-velocity spectrograph for \replaced{the 8.1\,m Gemini North telescope on Mauna Kea in Hawai'i}{one of the twin 6.5m Magellan Telescopes in Chile}, currently under construction at the University of Chicago \cite{maroonx}. While being developed for a particular instrument, it is suitable for calibrating any radial velocity spectrograph at the few cm\,s$^{-1}$ level after simple modifications to match the desired wavelength range and spectral resolution. Parts of the results have already been published in a SPIE proceeding \cite{stuermer2016}. In \S2 we describe the design of the FP etalon. In \S3 we describe our frequency stability monitoring of an etalon line using a scanning laser and rubidium gas cell. We present the performance of the system in \S4. We conclude with a discussion of applications and future development in \S5. \begin{figure*} \begin{center} \includegraphics[width=0.52\textwidth]{etalon_on_bb_annotated.pdf} \includegraphics[width=0.436\textwidth]{lesker2.pdf} \vspace{1mm} \caption{{FP etalon opto-mechanics and vacuum chamber}. Left: FP etalon (1) and OAP collimators (2) in vacuum compatible tip-tilt mounts (3) on a breadboard before vacuum integration. Right: System integrated in a vacuum chamber from J.K. Lesker with in-built channels for liquid circulation to provide temperature control at the \SI{\leq5}{\milli\K} level with an external bath thermostat (not shown). During operation the vacuum vessel is contained in another insulation box (also not shown here) to attenuate temperature variation of the room.} \label{lesker} \end{center} \end{figure*}
Our FP etalon calibrator combines for the first time the advantages of passively stabilized vacuum-gap bulk etalons \cite{wildi12} with the laser-locking technique demonstrated for single-mode fiber etalons \cite{gurevich14,schwab15}. By incorporating the advantages of illumination stability provided by a single mode fiber feed while avoiding the birefringence effects in SM etalons, we have created a calibrator that meets the needs of current and next generation RV spectrographs. The laser scanning technique allows us to use a FP etalon with moderate finesse for simultaneous comb creation and zero point tracking with high precision. Effective filtering of the laser light from the resulting etalon comb allows the usage of reference frequencies within the bandpass of the etalon comb without contamination of the white light comb or significant loss in comb coverage. We have demonstrated that we can measure the line position of our etalon with a precision of \SI{\leq3}{\cmps} relative to the rubidium reference during every exposure of a RV spectrograph. We can track the long-term drifts of the etalon with this precision at any time, removing any high-level intrinsic stability requirements on the etalon line position for timescales longer than a few seconds. In fact we can measure the \SI{\sim13}{\cmps} per day drift from aging of the etalon Zerodur spacer and apply a correction to the comb spectrum at the data reduction level. \deleted[]{ Due to the Zerodur aging alone, a passively stabilized etalon without referencing to rubidium would become useless as a primary wavelength calibrator after less than a day, if a stability of the calibration source at the $\SI{10}{\cmps}$ level is required.} The currently achieved level of precision is exceeding our requirements for MAROON-X, but we still see room for further improvement. As of this writing we have yet to thoroughly characterize the etalon as well as the rubidium reference in terms of its as-built behavior with respect to environmental changes compared to the theoretical expectations. This will ultimately improve the precision of the measured etalon line position but will also assure the long term absolute stability of our rubidium reference. For example, we find the stability of the mean laser power injected into the CoSy setup (\SI{\sim300}{\mu\watt}) to be just under \SI{1.7}{\%} P-V over the course of a week, which is already approaching our long term stability limit of \SI{5}{\mu\watt}. The laser power is currently not actively controlled and is thus subject to short term fluctuations and long drifts (aging of the diode). Further testing with the laser power modulated in a controlled fashion will reveal the true correlation between rubidium line frequency and laser power drift. Active control of the laser power or a post-processing correction for power induced shifts in our reference spectrum might become necessary. We also need to find the optimum scan rate to balance intrinsic (photon-limited) SNR and acoustic noise currently limiting the single scan precision. In order to apply the etalon spectrum for absolute wavelength calibration of a RV spectrograph, we need to determine the exact effective cavity length and the order number of each etalon line in the spectrum. An etalon with dielectric mirror coatings shows a wavelength dependent phase shift (a.k.a. group velocity dispersion) from the wavelength-dependent penetration depth of the light on the mirror surfaces. The amplitude is typically 100--300\,nm over the bandpass of the etalon, which translates into several \si{\kmps} deviation of the true comb line positions versus their ``ideal'' position, i.e. equally spaced in frequency. The group velocity dispersion of the etalon coatings is a known, smooth, and slowly varying function with wavelength \cite{wildi10, schwab15}. We can thus solve for the true line positions or compare the etalon spectrum with a ThAr reference spectrum \cite{bauer15}. Measuring and tracking our etalon at a single frequency might potentially impose a limitation on applying the etalon comb as a long-term primary wavelength calibrator for high-precision RV measurements. We are relying on the assumption that a drift measured for one etalon line is the same for all comb lines across the whole bandpass, modulo $\lambda_n/\lambda_{ref}$. This is a reasonable assumption, since after stabilizing the input illumination with a SM fiber, the most likely reason for wavelength-dependent effects are changes in the dielectric coatings of the FP etalon mirrors, affecting the penetration depth of the incoming light (and, hence, the cavity length $l$) in a time \textsl{and} wavelength dependent way. Aging and contamination of soft coatings is a known effect, although unlikely to occur for an etalon under vacuum. Nonetheless, minute changes in the coating thickness at the sub-nm level could cause \si{\mps} chromatic line shifts. Hence, simultaneous monitoring at a second reference wavelength or a temporary comparison to a LFC over a moderate amount of time is ultimately desirable to verify the absence of chromatic time-depended effects in a FP etalon.
16
7
1607.05172
We report on the construction and testing of a vacuum-gap Fabry-Pérot etalon calibrator for high precision radial velocity spectrographs. Our etalon is traced against a rubidium frequency standard to provide a cost effective, yet ultra precise wavelength reference. We describe here a turn-key system working at 500 to 900 nm, ready to be installed at any current and next-generation radial velocity spectrograph that requires calibration over a wide spectral bandpass. Where appropriate, we have used off-the-shelf, commercial components with demonstrated long-term performance to accelerate the development timescale of this instrument. Our system combines for the first time the advantages of passively stabilized etalons for optical and near-infrared wavelengths with the laser-locking technique demonstrated for single-mode fiber etalons. We realize uncertainties in the position of one etalon line at the 10 cm s<SUP>-1</SUP> level in individual measurements taken at 4 Hz. When binning the data over 10 s, we are able to trace the etalon line with a precision of better than 3 cm s<SUP>-1</SUP>. We present data obtained during a week of continuous operation where we detect (and correct for) the predicted, but previously unobserved shrinking of the etalon Zerodur spacer corresponding to a shift of 13 cm s<SUP>-1</SUP> per day.
false
[ "high precision radial velocity spectrographs", "ultra precise wavelength reference", "individual measurements", "single-mode fiber etalons", "s", "a wide spectral bandpass", "calibration", "one etalon line", "the etalon line", "day", "passively stabilized etalons", "demonstrated long-term performance", "Zerodur", "continuous operation", "Our etalon", "any current and next-generation radial velocity spectrograph", "the etalon Zerodur spacer", "a vacuum-gap Fabry-Pérot etalon calibrator", "10 cm s<", "a rubidium frequency standard" ]
13.706145
10.347553
11
12510981
[ "Heurtier, Lucien", "Teresi, Daniele" ]
2016PhRvD..94l5022H
[ "Dark matter and observable lepton flavor violation" ]
19
[ "Service de Physique Théorique—Université Libre de Bruxelles, Boulevard du Triomphe, CP225, 1050 Brussels, Belgium; Deutsches Elektronen-Synchrotron DESY, 22607 Hamburg, Germany", "Service de Physique Théorique—Université Libre de Bruxelles, Boulevard du Triomphe, CP225, 1050 Brussels, Belgium" ]
[ "2016JCAP...11..038K", "2016PhRvD..94g3004R", "2017IJMPA..3230023B", "2017IJMPA..3250078A", "2017JCAP...02..042H", "2017JCAP...05..027R", "2017JCAP...07..016C", "2017JHEP...05..102G", "2017PhRvD..95g5017B", "2017PhRvD..95k5037D", "2017PhRvD..96c5018H", "2017arXiv170303416R", "2017arXiv170702550D", "2017arXiv170900880A", "2018IJMPA..3342001D", "2018PhRvD..97g5022A", "2019PhRvD..99i5036K", "2019arXiv191208188M", "2021JHEP...01..169M" ]
[ "astronomy", "physics" ]
4
[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1975AnPhy..93..193F", "1977PhLB...67..421M", "1979NuPhB.155...52G", "1979PhRvL..42..407T", "1980PhRvD..22.2227S", "1980PhRvL..44..912M", "1981NuPhB.193..297G", "1981PhLB...98..265C", "1981PhLB...99..411G", "1986PhRvD..34.1642M", "1986PhRvL..56..561M", "1988PhRvD..38..448H", "1992ZPhyC..55..275P", "1993PhLB..318..360B", "1994PhRvD..49.2398P", "1995NuPhB.437..491I", "1997PhRvD..56.5431P", "2000PhRvD..62b3004K", "2002PhRvL..88i1304M", "2003PhRvD..67g3015F", "2004NuPhB.692..303P", "2004PhRvL..92j1601B", "2005PhLB..631..151A", "2005PhRvL..95h1602P", "2005PhRvL..95p1801M", "2006JHEP...06..053A", "2006PhLB..639..414S", "2006PhRvL..97x1301K", "2007PhRvD..75i4001L", "2008PhRvD..77f5014P", "2008PhRvD..78a3008P", "2009JHEP...09..038G", "2010JHEP...03..080H", "2010JHEP...04..023B", "2010PhLB..690..145G", "2011JCAP...03..042A", "2011PhRvX...1b1026F", "2012EPJC...72.2058L", "2012JCAP...03..029B", "2012JCAP...05..034C", "2012JHEP...06..031E", "2012PhLB..709..222E", "2012arXiv1211.7019M", "2013ApJS..208...19H", "2013JHEP...01..118A", "2013PhLB..719..415L", "2013PhRvD..88d3502V", "2013PhRvD..88i3010L", "2014ChPhC..38i0001O", "2014JCAP...01..034C", "2014NuPhB.886..569D", "2014PhLB..735...69Q", "2014PhLB..736..494B", "2015JCAP...06..011M", "2015JCAP...07..014F", "2015JHEP...05..053A", "2015JHEP...10..067H", "2015NJPh...17g5019D", "2015NuPhB.897..749B", "2015PhRvD..91i5010M", "2016A&A...594A..13P", "2016JHEP...08..157H", "2016JHEP...12..007A", "2016JHEP...12..150D", "2016NPPP..273.1883B", "2016PhRvD..93k1703M", "2016PhRvD..94h5012F", "2016PhRvD..94j3009L", "2016PhRvL.117i1801H", "2016RPPh...79l4201A", "2017JCAP...01..025A", "2017JCAP...02..042H" ]
[ "10.1103/PhysRevD.94.125022", "10.48550/arXiv.1607.01798" ]
1607
1607.01798_arXiv.txt
Among other problems, the Standard Model (SM) of particle physics lacks an explanation of what is the dark constituent of our Universe, as well as the origin of the tiny neutrino masses. As for the latter, the most popular paradigm is to extend the SM with additional fermions, namely right-handed (RH) neutrinos, whose role is to generate tiny masses for the active neutrinos, via the so called seesaw mechanism~\cite{Seesaw}. At the same time, one of the RH neutrinos can play the role of dark matter (DM). The simplest formulation of this is in the type-I seesaw scenario~\cite{Seesaw}, where two of the RH neutrinos are responsible for the active-neutrino masses and mixing, whereas the third one can play the role of warm DM~\cite{Asaka:2005an,keVreview}. The Majorana masses for the RH neutrinos $N_R$ can, in turn, be generated by the spontaneous breaking of a global $U(1)$ symmetry~\cite{Majoron}. In this so-called \emph{Majoron model} an additional complex scalar field is added to the theory to break $U(1)_l$, thus generating Majorana masses for the RH neutrinos, also entailing an interesting phenomenology coming from the presence of the scalar and pseudo-scalar couplings to both active and sterile neutrinos~\cite{Majoron_pheno,ApostolosMajoron}. A drawback of the ``vanilla'' seesaw model is that the smallness of the active-neutrino masses requires the Yukawa couplings of the RH neutrinos to be very small for RH-neutrino masses at reach of current or near-future experiments, thus rendering the model difficult to test in the foreseeable future. The required Yukawa couplings are of order \begin{equation}\label{eq:naive_seesaw} h_{seesaw} \ \approx \ 6 \, \times \, 10^{-8} \, \times \, \sqrt{\frac{m_{\nu_L}}{0.1 \,\mathrm{eV}}} \, \times \, \sqrt{\frac{m_{N_R}}{\mathrm{GeV}}}\;. \end{equation} However, a number of variants of the type-I seesaw mechanism have been developed (e.g the inverse seesaw~\cite{inverseseesaw}, linear seesaw~\cite{linearseesaw}, etc.), where the presence of a leptonic symmetry $U(1)_l$ protects the smallness of the active neutrino masses, thus allowing for much larger Yukawa couplings\footnote{This can be achieved also by means of a discrete symmetry, see e.g.~\cite{Dev:2013oxa}.} than in \eqref{eq:naive_seesaw}, even of order $10^{-3}$ or higher~\cite{largeyukawas}. Therefore, these models provide a way to test the seesaw mechanism in the near future, for instance by the observation of lepton-flavour-violation (LFV) processes, such as $\mu \to e \gamma$ and $\mu \to e$ conversion in nuclei. In particular, the sensitivity of the latter will improve by several orders of magnitude in the near future, thanks to the planned experiments Mu2e and COMET, as well as to the more distant proposal PRISM/PRIME. Typically, in this class of models, in order to generate the observed pattern of neutrino masses and mixing, the leptonic symmetry is explicitly broken \emph{by hand} in different possible ways, giving rise to the so-called inverse-seesaw or linear-seesaw textures, for instance. Since this class of leptonic symmetries generically gives two quasi-degenerate RH neutrinos, it is tempting to try to explain also the Baryon Asymmetry of the Universe in this model, via the resonant leptogenesis mechanism~\cite{rl, Dev:2014laa}. This can be achieved by supplementing the leptonic symmetry with a larger $O(3)$ symmetry in the RH sector~\cite{Dev:2014laa}. However, it appears to be difficult to reconcile observable LFV rates and successful leptogenesis in the \emph{minimal} models possessing only the leptonic symmetry $U(1)_l$ (see~\cite{Blanchet:2009kk} and Appendix~A of~\cite{Dev:2014laa}). This is true even if one considers GeV-scale leptogenesis mechanisms via RH-neutrino oscillations (see e.g.~\cite{ARSpheno}) or Higgs lepton-number violating decays~\cite{Hambye:2016sby}. Therefore, is it natural to try to address an alternative question: \emph{is it instead possible, in this class of models, to have observable LFV rates and a successful DM candidate?} In this paper we construct a model achieving this, in which the same leptonic symmetry $U(1)_l$ responsible for (i) light-neutrino masses with large Yukawa couplings, at the same time (ii) stabilizes one of the RH neutrinos, which is therefore a DM candidate. The spontaneous breaking of $U(1)_l$ involves a \emph{Majoron-like} complex scalar field, which in turn (iii) provides a mechanism to generate successfully the DM candidate in the early Universe, together with its keV-scale mass. The charge assignment under $U(1)_l$ needed to achieve this gives, at the same time, a particular pattern for the breaking of the leptonic symmetry, which in our model is \emph{not} performed by hand, but is instead related to the above points. After this introduction, in section~\ref{sec:model} we will construct the model, derive the mass matrix of the neutrinos after the spontaneous breaking of the electroweak symmetry and $U(1)_l$, and describe quantitatively the generation of DM. In section~\ref{sec:pheno} we study the phenomenology of the model, in particular at near-future $\mu \to e$ conversion experiments, as well as in direct searches of the RH neutrinos. We also study the interactions of the Majoron-like field, which can give detectable imprints at the observation of future supernovae. Finally, in section~\ref{sec:conclusion} we draw our conclusions.
\label{sec:conclusion} In this paper, we have investigated the possibility that a leptonic global symmetry $U(1)_l$ protecting the lightness of the active neutrinos in the type-I seesaw scenario, at the same time allowing for sizeable Yukawa couplings with the RH sector, is broken spontaneously by adding one ``Majoron-like'' complex scalar field to the seesaw Lagrangian. The charge assignment under $U(1)_l$ of such setup allows to render one of the RH neutrino absolutely stable, which therefore constitutes a natural DM candidate. Whereas in this framework it would be difficult to produce the DM candidate via its (tiny) interactions with the active sector, the decays of the field responsible for the breaking of $U(1)_l$ allow for a successful production of DM in the early Universe, via a freeze-in mechanism. In the dimension-4 Lagrangian only 3 ``flavoured'' couplings are present. In order to obtain a sufficient number of couplings to fit the non-trivial active-neutrino mass matrix, dimension-5 effective operators need to be considered too. These can arise from some $U(1)_l$ invariant heavy sector at an intermediate scale $\Lambda$. Interestingly, the requirement of having large LFV rates, observable in the near future, as well as the observed DM relic density, fixes the various scales of the model: \begin{itemize}\itemsep0.3em \item[$(i)$] the DM mass has to be in $\mathrm{keV}$ range, with possibilities up to the $\mathrm{MeV}$ range, in some regions of the parameter space; \item[$(ii)$]the scale $\Lambda$ has to be at least $10^{13}\, \mathrm{GeV}$ for $c_h \simeq 1$, which coincides with the scale of intermediate breaking of various GUT models, or with the GUT scale itself; \item[$(iii)$] the scale of $U(1)_l$ breaking is typically in the $10-1000\,\mathrm{TeV}$ range. As pointed out in~\cite{ApostolosMajoron}, the phase transition breaking $U(1)_l$ can even coincide with the electroweak one. \end{itemize} The model -- in addition to be as minimal as possible -- has a set of features which make it testable by future neutrino-physics measurements. A first point is that the requirement of large LFV rates is more easily satisfied for an inverted-hierarchy mass spectrum, although there are possibilities even for a normal-hierarchy spectrum too. More importantly, since only two RH neutrinos have an active role in the seesaw mechanism, the lightest of the active neutrinos is automatically massless. By construction, the presence of large Yukawa couplings allows for LFV processes with large rates, detectable in the near-future at $\mu \to e$ conversion experiments Mu2e and COMET. As we have explained above, this requirement, which is the original motivation for the model, fixes its mass scales. In addition to this, since the heavier RH states have masses lighter than 300 GeV in a good portion of the parameter space, the model can be tested also by the direct production of these states at future proposed experiments, such as SHiP, FCC-ee and ILC. Notice that this feature is precisely what distinguishes the model from a ``standard'' Majoron one, since the latter cannot yield a sufficiently large mixing with the SM leptons, unless a severe fine-tuning in the seesaw relation is invoked. Notice also that, along the same lines of the model discussed here, one could construct models with a larger number of fields and where the neutrino mass generation is, for instance, only of the linear-seesaw or inverse-seesaw type, cf.~\eqref{eq:cont1} and \eqref{eq:cont2}. Their phenomenology would be similar to the one discussed here, apart from the number of states in the GeV-TeV range, which needs to be larger in the inverse-seesaw case, and the suppression of the couplings of the scalar to SM leptons, in the linear-seesaw case. Finally, the coupling of the Majoron-like scalar field to the active neutrinos can be large in a region of the parameter space with $\Lambda \sim 10^{13-14}\, \text{GeV}$, being in particular close to the recent bound obtained from the neutrino burst of supernovae. Therefore, the model can have addition complementary signatures at future supernova detections by IceCube and SuperKamiokande, which would provide an additional strong piece of evidence for it. \medskip
16
7
1607.01798
Seesaw models with leptonic symmetries allow right-handed (RH) neutrino masses at the electroweak scale, or even lower, at the same time having large Yukawa couplings with the Standard Model leptons, thus yielding observable effects at current or near-future lepton-flavor-violation (LFV) experiments. These models have been previously considered also in connection to low-scale leptogenesis, but the combination of observable LFV and successful leptogenesis has appeared to be difficult to achieve unless the leptonic symmetry is embedded into a larger one. In this paper, instead, we follow a different route and consider a possible connection between large LFV rates and dark matter (DM). We present a model in which the same leptonic symmetry responsible for the large Yukawa couplings guarantees the stability of the DM candidate, identified as the lightest of the RH neutrinos. The spontaneous breaking of this symmetry, caused by a Majoron-like field, also provides a mechanism to produce the observed relic density via the decays of the latter. The phenomenological implications of the model are discussed, finding that large LFV rates, observable in the near-future μ →e conversion experiments, require the DM mass to be in the keV range. Moreover, the active-neutrino coupling to the Majoron-like scalar field could be probed in future detections of supernova neutrino bursts.
false
[ "large LFV rates", "large Yukawa couplings", "observable LFV", "supernova neutrino bursts", "DM", "observable effects", "LFV", "leptonic symmetries", "successful leptogenesis", "future detections", "dark matter", "the large Yukawa couplings", "Standard Model", "Majoron", "leptonic", "Yukawa", "the DM mass", "the keV range", "the DM candidate", "the RH neutrinos" ]
8.618865
-1.760809
54
2074426
[ "Grayson, J. A.", "Ade, P. A. R.", "Ahmed, Z.", "Alexander, K. D.", "Amiri, M.", "Barkats, D.", "Benton, S. J.", "Bischoff, C. A.", "Bock, J. J.", "Boenish, H.", "Bowens-Rubin, R.", "Buder, I.", "Bullock, E.", "Buza, V.", "Connors, J.", "Filippini, J. P.", "Fliescher, S.", "Halpern, M.", "Harrison, S.", "Hilton, G. C.", "Hristov, V. V.", "Hui, H.", "Irwin, K. D.", "Kang, J.", "Karkare, K. S.", "Karpel, E.", "Kefeli, S.", "Kernasovskiy, S. A.", "Kovac, J. M.", "Kuo, C. L.", "Leitch, E. M.", "Lueker, M.", "Megerian, K. G.", "Monticue, V.", "Namikawa, T.", "Netterfield, C. B.", "Nguyen, H. T.", "O'Brient, R.", "Ogburn, R. W.", "Pryke, C.", "Reintsema, C. D.", "Richter, S.", "Schwarz, R.", "Sorenson, C.", "Sheehy, C. D.", "Staniszewski, Z. K.", "Steinbach, B.", "Teply, G. P.", "Thompson, K. L.", "Tolan, J. E.", "Tucker, C.", "Turner, A. D.", "Vieregg, A. G.", "Wandui, A.", "Weber, A. C.", "Wiebe, D. V.", "Willmert, J.", "Wu, W. L. K.", "Yoon, K. W." ]
2016SPIE.9914E..0SG
[ "BICEP3 performance overview and planned Keck Array upgrade" ]
51
[ "Stanford Univ. (United States)", "Cardiff Univ. (United Kingdom)", "Kavli Institute for Particle Astrophysics &amp; Cosmology (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "The Univ. of British Columbia (Canada)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Univ. of Toronto (Canada)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "California Institute of Technology (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Univ. of Minnesota, Minneapolis (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "California Institute of Technology (United States)", "Univ. of Minnesota, Minneapolis (United States)", "The Univ. of British Columbia (Canada)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "National Institute of Standards and Technology (United States)", "California Institute of Technology (United States)", "California Institute of Technology (United States)", "Stanford Univ. (United States)", "Stanford Univ. (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Stanford Univ. (United States)", "California Institute of Technology (United States)", "Stanford Univ. (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Stanford Univ. (United States)", "Kavli Institute for Cosmological Physics, The Univ. of Chicago (United States)", "California Institute of Technology (United States)", "Jet Propulsion Lab. (United States)", "Stanford Univ. (United States)", "Stanford Univ. (United States)", "Univ. of Toronto (Canada)", "Jet Propulsion Lab. (United States)", "California Institute of Technology (United States)", "Stanford Univ. (United States)", "Univ. of Minnesota (United States)", "National Institute of Standards and Technology (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Univ. of Minnesota (United States)", "Harvard-Smithsonian Ctr. for Astrophysics (United States)", "Univ. of Minnesota, Minneapolis (United States)", "California Institute of Technology (United States)", "California Institute of Technology (United States)", "Univ. of California, San Diego (United States)", "Stanford Univ. (United States)", "Stanford Univ. (United States)", "Cardiff Univ. (United Kingdom)", "Jet Propulsion Lab. (United States)", "Harvard-Smithsonian Center for Astrophysics (United States)", "Stanford Univ. (United States)", "Jet Propulsion Lab. (United States)", "Univ. of British Columbia (Canada)", "Univ. of Minnesota (United States)", "Univ. of California, Berkeley (United States)", "Stanford Univ. (United States)" ]
[ "2016arXiv160803707Z", "2017PhRvD..95d3523N", "2017PhRvD..96l3521G", "2017arXiv170602464A", "2018ApPhL.112m2601C", "2018JCAP...02..009B", "2018JCAP...04..016F", "2018JCAP...08..038R", "2018JLTP..193..562S", "2018JLTP..193..633H", "2018JLTP..193..876C", "2018JLTP..193..976N", "2018PhR...761....1B", "2018PhRvD..97f3505N", "2018PhRvD..98l3515K", "2018SPIE10708E..19H", "2018SPIE10708E..2KB", "2018arXiv180510475A", "2019BAAS...51c.567G", "2019MNRAS.484.1616Z", "2019PhRvD..99b3530N", "2019PhRvD..99d3529E", "2020ApOpt..59.5439E", "2020JHEP...04..042H", "2020JHEP...05..155H", "2020JLTP..200..342B", "2020PASJ...72....6N", "2020PDU....2700450R", "2020PhRvD.102h3530G", "2020PhRvD.102j3005C", "2020SPIE11453E..14M", "2021JCAP...08..008A", "2021PhLB..81536137G", "2021PhRvD.103j3531C", "2021PhRvD.104j3512C", "2022ApJ...927...77A", "2022ApJS..262...52A", "2022EPJC...82.1146P", "2022JCAP...01..039K", "2022PhRvD.106f3505G", "2022arXiv220305010E", "2023ApOpt..62.4334G", "2023JCAP...03..061C", "2023JCAP...09..019G", "2023JCAP...09..042D", "2023MNRAS.520.1757G", "2023SuScT..36k5004L", "2023Univ....9..302H", "2023arXiv231103199O", "2024arXiv240201233Z", "2024arXiv240412779G" ]
[ "astronomy", "physics" ]
6
[ "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "2003RScI...74.3807D", "2006SPIE.6275E..0UA", "2008JLTP..151..908B", "2012SPIE.8452E..1GO", "2014ApJ...783...67B", "2014ApJ...794..171P", "2014JLTP..176..835A", "2014PhRvL.112x1101B", "2014SPIE.9153E..1NA", "2014SPIE.9153E..3BK", "2015ApJ...807..151K", "2015ApJ...812..176B", "2015PhDT.......278W", "2015PhRvL.114j1301B", "2016PhRvL.116c1302B", "2016SPIE.9914E..30K" ]
[ "10.1117/12.2233894", "10.48550/arXiv.1607.04668" ]
1607
1607.04668_arXiv.txt
16
7
1607.04668
Bicep3 is a 520mm aperture, compact two-lens refractor designed to observe the polarization of the cosmic microwave background (CMB) at 95 GHz. Its focal plane consists of modularized tiles of antenna-coupled transition edge sensors (TESs), similar to those used in Bicep2 and the Keck Array. The increased per-receiver optical throughput compared to Bicep2/Keck Array, due to both its faster f=1:7 optics and the larger aperture, more than doubles the combined mapping speed of the Bicep/Keck program. The Bicep3 receiver was recently upgraded to a full complement of 20 tiles of detectors (2560 TESs) and is now beginning its second year of observation (and first science season) at the South Pole. We report on its current performance and observing plans. Given its high per-receiver throughput while maintaining the advantages of a compact design, Bicep3- class receivers are ideally suited as building blocks for a 3rd-generation CMB experiment, consisting of multiple receivers spanning 35 GHz to 270 GHz with total detector count in the tens of thousands. We present plans for such an array, the new "BICEP Array" that will replace the Keck Array at the South Pole, including design optimization, frequency coverage, and deployment/observing strategies.
false
[ "Keck Array", "total detector count", "design optimization", "BICEP Array", "Bicep3- class receivers", "frequency coverage", "multiple receivers", "thousands", "first science season", "detectors", "Bicep2/Keck Array", "receiver", "the Keck Array", "deployment/observing strategies", "CMB", "modularized tiles", "Bicep2", "the South Pole", "Bicep/Keck", "observation" ]
10.636029
2.810985
68
12409356
[ "Vican, Laura", "Schneider, Adam", "Bryden, Geoff", "Melis, Carl", "Zuckerman, B.", "Rhee, Joseph", "Song, Inseok" ]
2016ApJ...833..263V
[ "Herschel Observations of Dusty Debris Disks" ]
23
[ "Department of Physics and Astronomy, University of California, Los Angeles, CA 90095, USA", "Department of Physics and Astronomy, University of Toledo, Toledo, OH 43606, USA", "JPL/Caltech, Pasadena, CA 01109, USA", "Center for Astrophysics and Space Sciences, University of California, San Diego, CA 92093-0424, USA", "Department of Physics and Astronomy, University of California, Los Angeles, CA 90095, USA", "California State Polytechnic University Pomona, 3801 West Temple Avenue, Pomona, CA 91768", "Department of Physics and Astronomy, University of Georgia, Athens, GA 30602-2451, USA" ]
[ "2017ApJ...840L..20L", "2017ApJ...845..120B", "2017ApJ...849..123M", "2017MNRAS.470.3606H", "2018ApJ...854...53C", "2018exha.book.....P", "2019ApJ...875...45T", "2019MmSAI..90..543M", "2020A&A...639A..11R", "2020MNRAS.497.2811K", "2020PASP..132j2001H", "2021ApJ...910...27M", "2021ApJ...922...75S", "2021ApJ...923...90M", "2021MNRAS.501.6168M", "2021RAA....21...60L", "2022AJ....163...25L", "2023A&A...671L...2B", "2023ApJ...951..111C", "2023MNRAS.520.3218M", "2023MNRAS.521.5940M", "2024MNRAS.528.4528M", "2024MNRAS.531.4482B" ]
[ "astronomy" ]
5
[ "circumstellar matter", "infrared: planetary systems", "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1975Icar...24..504H", "1984ApJ...278L..23A", "1995SPIE.2475..360B", "1999ApJ...512..377H", "2001ApJ...560L.171C", "2003ApJ...598..626D", "2003ApJ...599..342S", "2004ApJ...603..738Z", "2004ApJS..154....1W", "2004ApJS..154...25R", "2004SPIE.5487...62H", "2005AJ....130..269K", "2005ApJ...623..493C", "2005Natur.436..363S", "2007ApJ...658..569W", "2007ApJ...660.1556R", "2007ApJS..173..143B", "2008ARA&A..46..339W", "2008ApJ...673.1106L", "2008ApJ...675..777R", "2008ApJ...688.1345Z", "2008ApJS..179..451K", "2008SPIE.7010E..04D", "2008SPIE.7010E..05P", "2008SPIE.7010E..06G", "2008Sci...322.1348M", "2009A&A...508..833D", "2009ApJ...698.1068P", "2009ApJ...701.1367C", "2009ApJ...701.2019L", "2009ApJ...705.1226B", "2010A&A...518L...1P", "2010A&A...518L.135M", "2010AJ....140.1868W", "2010ASPC..434..139O", "2010ApJ...708.1611R", "2010ApJ...717L..57M", "2010Natur.468.1080M", "2011A&A...530A.138C", "2011ARA&A..49...67W", "2011ApJ...726...72W", "2011ApJ...730L..29M", "2011ApJ...732...61Z", "2011MNRAS.410..190T", "2012A&A...537A.120Z", "2012A&A...542A..90O", "2012ApJ...745..147R", "2012ApJ...748...30R", "2012ApJ...749L..29F", "2012ApJ...752...58Z", "2012ApJ...758...77Z", "2012MNRAS.425..657J", "2012MNRAS.427..343M", "2013A&A...551A.134O", "2013A&A...553A.132M", "2013ApJ...775...55B", "2013ApJ...776...77K", "2013ApJ...776..111M", "2013ApJ...777..160B", "2013ApJ...778...12M", "2013ApJ...779L..26R", "2013MNRAS.428.1263B", "2014AJ....147..146K", "2014ApJ...784...33G", "2014ApJ...792...65P", "2014ApJS..211...25C", "2014MNRAS.444.3164K", "2014PASP..126..827H", "2015ApJ...798...86Z", "2015ApJ...798...87M", "2015ApJ...799..146S", "2015ApJ...801..143M", "2015MNRAS.447..577M", "2015MNRAS.449.3160R", "2016ApJ...826...19N", "2016MNRAS.459.2893M" ]
[ "10.3847/1538-4357/833/2/263", "10.48550/arXiv.1607.03754" ]
1607
1607.03754_arXiv.txt
Debris disks are signposts of planetesimal formation and, as such, are crucial subjects for study when considering the evolution of a planetary system. Debris disks are typically identified via an excess infrared (IR) flux above the stellar photosphere. New debris disks have been discovered with five satellites, starting with the \textit{InfraRed Astronomical Satellite (IRAS)} discovery of the first debris disk around Vega in 1984 \citep{Aumann_1984}. In 1995, the Infrared Space Observatory (ISO; \citealt{Boulade_1995}) imaged the sky at wavelengths ranging from 2.5 to 240 $\mu$m, providing low-resolution long wavelength photometry. Subsequently, the \textit{Spitzer Space Telescope} \citep{Werner_2004} provided mid-IR spectroscopy with the Infrared Spectrograph (IRS; \citealt{Houck_2004}), and photometry with the Multiband Imaging Photometer (MIPS; \citealt{Rieke_2004}), leading to the detection of over 100 new debris disks (e.g. \citealt{Zuckerman_2004}, \citealt{Plavchan_2009}, \citealt{Chen_2009}, \citealt{Chen_2014}). The Wide Field Survey Explorer (WISE; \citealt{Wright_2010}) imaged the sky at 3.4, 4.6, 11, and 22$\mu$m, and was also most sensitive to warm and hot debris disks that peak in the mid-IR. In 2009, the \textit{Herschel Space Observatory} \citep{Pilbratt_2010} began taking data, providing far-IR photometry with the Photodetector Array Camera and Spectrometer (PACS; \citealt{Poglitsch_2008}) and spectroscopy with the Spectral and Photometric Imaging Receiver (SPIRE; \citealt{Griffin_2008}) and the Heterodyne Instrument for the Far Infrared (HIFI; \citealt{deGraauw_2008}). Where Spitzer was uniquely able to detect warm ($>$100 K) debris disks in the terrestrial planet zone, Herschel was sensitive to cooler disks ($<$100 K) at larger radial separations from their host stars. In this paper, we present Herschel observations of 24 stars initially identified with IRAS and/or Spitzer as definitely or possibly possessing a debris disk. One goal of our project was to search for cold dust components that would peak near the Herschel PACS wavelengths (70, 100, and 160 $\mu$m). In the case that there is no separate cold dust component present, Herschel photometry helps to better characterize the Rayleigh-Jeans tail of thermal emission from warm dust. Another goal was the identification of disks with double-belt debris systems; that is, systems containing an inner belt of warm or hot ($>$100K) dust and an outer belt of cold ($<$100 K) dust. Such systems would be a direct analog of our own Solar System, which hosts a Kuiper Belt that lies between 30-50 AU at $\sim$50 K and an Asteroid Belt at 3 AU and $\sim$175 K. A double-belt system may also be a signature of a planet (or planets) that lie in the gap between dust belts. Such systems have been discovered around HR 8799 (\citealt{Marois_2008}, \citealt{Marois_2010}, \citealt{Matthews_2010}) and HD 95086 (\citealt{Rameau_2013}, \citealt{Su_2015}).
We observed 24 stars with the PACS camera on the Herschel Space Observatory. Two infrared sources detected with Herschel are offset from the target coordinates by $>$2$\arcsec$. These are unlikely to be due to dusty debris belts. Three targets were not detected with Herschel, and are non-excess stars. One star was detected with Herschel, but shows no evidence of an IR excess. Two stars have low IR fluxes, and could possibly be explained by contamination by a background galaxy at z$\sim$1. One target star is clearly contaminated by a background galaxy. The remaining 15 stars were examined to determine dust properties and the possibility of ongoing planet formation. A summary of the results can be found in Table 5. \begin{itemize} \item Nine stars (HD 15407, HD 23514, HD 35650, HD 54341, HD 76543, HD 76582, HD 113766, HD 121191, and HD 131488) appear to have dust components at two temperatures according to SED fitting. All appear to have spatially separated dust belts. \item Three stars (HD 15407, HD 113766, and BD+20 307) have disks that cannot be explained by a steady-state collisional process alone. Two other stars (HD 23514 and HD 124718) are likely explained by a transient process. The other debris disks could be explained by steady-state collisional processes. \item Dust belts at HD 54341, HD 84870, HD 85672, and HD 121191 are at large enough distances from their host stars that the dust could be stirred by a yet-unseen planet or binary companion. Stirring of these disks by putative 1000 km-sized bodies may be insufficient to explain the dust production in these systems. \item The stars HD 54341, HD 76543, HD 76582, HD 84870, HD 85672, and HD 99945 are spatially resolved (or probably resolved as in the case of HD 121191), and a comparison between their blackbody radii (from SED-fitting) and resolved radii show that the latter are typically larger than the former, perhaps by up to a factor as large as 10 (HD 121191). \item Six stars (HD 15407, HD 23514, HD113766, HD 121191, HD 124718, and BD+20 307) show no evidence of cold ($<$100 K) dust in the Herschel data. \end{itemize}
16
7
1607.03754
We present results from two Herschel observing programs using the Photodetector Array Camera and Spectrometer. During three separate campaigns, we obtained Herschel data for 24 stars at 70, 100, and 160 μm. We chose stars that were already known or suspected to have circumstellar dust based on excess infrared (IR) emission previously measured with the InfraRed Astronomical Satellite (IRAS) or Spitzer and used Herschel to examine long-wavelength properties of the dust. Fifteen stars were found to be uncontaminated by background sources and possess IR emission most likely due to a circumstellar debris disk. We analyzed the properties of these debris disks to better understand the physical mechanisms responsible for dust production and removal. Seven targets were spatially resolved in the Herschel images. Based on fits to their spectral energy distributions, nine disks appear to have two temperature components. Of these nine, in three cases, the warmer dust component is likely the result of a transient process rather than a steady-state collisional cascade. The dust belts at four stars are likely stirred by an unseen planet and merit further investigation.
false
[ "circumstellar dust", "dust production", "merit further investigation", "Herschel data", "IR emission", "Spectrometer", "Herschel", "removal", "results", "the warmer dust component", "observing programs", "a circumstellar debris disk", "stars", "Spitzer", "IRAS", "background sources", "a steady-state collisional cascade", "IR", "The dust belts", "these debris disks" ]
9.160601
13.007449
144
4928065
[ "Petrushevska, T.", "Amanullah, R.", "Goobar, A.", "Fabbro, S.", "Johansson, J.", "Kjellsson, T.", "Lidman, C.", "Paech, K.", "Richard, J.", "Dahle, H.", "Ferretti, R.", "Kneib, J. P.", "Limousin, M.", "Nordin, J.", "Stanishev, V." ]
2016A&A...594A..54P
[ "High-redshift supernova rates measured with the gravitational telescope A 1689" ]
31
[ "Oskar Klein Centre, Physics Department, Stockholm University, 106 91, Stockholm, Sweden ; Physics Department, Stockholm University, 106 91, Stockholm, Sweden", "Oskar Klein Centre, Physics Department, Stockholm University, 106 91, Stockholm, Sweden; Physics Department, Stockholm University, 106 91, Stockholm, Sweden", "Oskar Klein Centre, Physics Department, Stockholm University, 106 91, Stockholm, Sweden; Physics Department, Stockholm University, 106 91, Stockholm, Sweden", "NRC Herzberg Institute for Astrophysics, 5071 West Saanich Road, Victoria V9E 2E7, British Columbia, Canada", "Department of Particle Physics and Astrophysics, Weizmann Institute of Science, 7610001, Rehovot, Israel", "Physics Department, Stockholm University, 106 91, Stockholm, Sweden", "Australian Astronomical Observatory, PO Box 915, 1670, North Ryde, Australia", "Universitäts-Sternwarte, Fakultät für Physik, Ludwig-Maximilians Universitaet München, Scheinerstr. 1, 81679, Muenchen, Germany; Excellence Cluster Universe, Boltzmannstr. 2, 85748, Garching, Germany", "Univ. Lyon, Univ. Lyon1, Ens de Lyon, CNRS, Centre de Recherche Astrophysique de Lyon, UMR 5574, 69230, Saint-Genis-Laval, France", "Institute of Theoretical Astrophysics, University of Oslo, PO Box 1029, Blindern, 0315, Oslo, Norway", "Oskar Klein Centre, Physics Department, Stockholm University, 106 91, Stockholm, Sweden; Physics Department, Stockholm University, 106 91, Stockholm, Sweden", "Laboratoire d'Astrophysique, École Polytechnique Fédérale de Lausanne (EPFL), Observatoire de Sauverny, 1290, Versoix, Switzerland", "Laboratoire d'Astrophysique de Marseille, UMR 6610, CNRS-Université de Provence, 38 rue Frédéric Joliot-Curie, 13388, Marseille Cedex 13, France", "Institut für Physik, Humboldt-Universitat zu Berlin, Newtonstr. 15, 12489, Berlin, Germany", "Department of Physics, Chemistry and Biology, IFM, Linköping University, 581 83, Linköping, Sweden" ]
[ "2017ApJ...834L...5G", "2017MNRAS.472..616S", "2018A&A...614A.103P", "2018ARep...62..917P", "2018ApJ...864...91S", "2018ApJ...866...65R", "2018JCAP...06..019J", "2019ApJS..243....6G", "2019MNRAS.482..870E", "2019MNRAS.488.5300C", "2019PDU....2600397Y", "2019RPPh...82l6901O", "2020MNRAS.491.2447R", "2020MNRAS.496.3270M", "2020MNRAS.497.2201S", "2020PhRvD.101j3017B", "2020Symm...12.1966P", "2021MNRAS.504.5621D", "2022ApJS..259...58L", "2022ChPhL..39k9801L", "2022MNRAS.514.1315B", "2022MNRAS.517.2471Z", "2022arXiv220709632S", "2023A&A...680L...9A", "2023ApJ...948..115P", "2023MNRAS.518.3475T", "2023MNRAS.522.4718G", "2023MNRAS.525..542M", "2023MNRAS.526.5911J", "2024SSRv..220...13S", "2024arXiv240518589T" ]
[ "astronomy" ]
5
[ "supernovae: general", "gravitational lensing: strong", "galaxies: star formation", "galaxies: clusters: individual: A 1689", "techniques: photometric", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1937ApJ....86..217Z", "1938ApJ....88..529Z", "1955ApJ...121..161S", "1964MNRAS.128..307R", "1979ApJ...232...18H", "1988ApJ...332...75N", "1988ApJ...335L...9K", "1989ApJ...345..245C", "1996A&AS..117..393B", "1998ApJ...502..177H", "1998ApJ...503..325A", "1998MNRAS.296..763K", "1999A&A...351..459C", "2000A&A...363..476B", "2000A&AS..144..363A", "2000ApJ...532..679P", "2000ApJ...533..682C", "2000MNRAS.318.1093F", "2000MNRAS.319..549S", "2001ApJ...556L..71H", "2002AJ....123..745R", "2002ASPC..281..228B", "2002ApJ...573..144D", "2002MNRAS.332...37G", "2003A&A...405..859G", "2003ApJ...592...17B", "2003ApJS..149..289B", "2003MNRAS.338L..25O", "2004ApJ...613..189D", "2004SPIE.5492.1763P", "2004ogci.conf..183C", "2005A&A...437..789N", "2005ApJ...621...53B", "2005ApJ...629L..85S", "2005MNRAS.362..671R", "2006A&A...447...31A", "2006AJ....132.1126N", "2006ASPC..351..112B", "2006SSRv..123..485G", "2007ApJ...660....1O", "2007ApJ...660.1165S", "2007ApJ...663.1187H", "2007ApJ...664..761M", "2007ApJ...666.1024B", "2007ApJ...668..643L", "2007MNRAS.377.1229M", "2007NJPh....9..447J", "2008A&A...479...49B", "2008A&A...485..633L", "2008A&A...486..375A", "2008AJ....135.1343G", "2008ApJ...676..767T", "2008ApJ...681..462D", "2008MNRAS.383.1121M", "2009A&A...499..653B", "2009A&A...507...61S", "2009A&A...507...71G", "2009ApJ...695.1334S", "2009ApJ...701.1336L", "2009JCAP...01..047L", "2009arXiv0912.0201L", "2010ApJ...713.1026D", "2010ApJ...713.1073S", "2010ApJ...718..876S", "2010ApJ...722.1879M", "2010ApJ...723.1678C", "2011A&A...536A..94R", "2011AJ....141...19B", "2011ARNPS..61..251G", "2011ApJ...729..143C", "2011ApJ...738..154H", "2011ApJ...742L...7A", "2011CBET.2642....1A", "2011MNRAS.412.1441L", "2011MNRAS.417..916G", "2012A&A...537A.132B", "2012A&A...545A..96M", "2012ApJ...745...32B", "2012ApJ...754...17F", "2012ApJ...756..111M", "2012ApJ...757...70D", "2012ApJS..199...25P", "2013A&A...555A..10T", "2013MNRAS.435L..33P", "2013arXiv1305.5422S", "2014AJ....147..118R", "2014AJ....148...13R", "2014ARA&A..52..415M", "2014ApJ...780..143A", "2014ApJ...786....9P", "2014ApJ...786...67A", "2014ApJ...789...51Z", "2014ApJ...792..135T", "2014MNRAS.440.2742N", "2014htu..conf...27B", "2015A&A...584A..62C", "2015ApJ...806..207U", "2015ApJ...811...70R", "2015ApJ...812..114T", "2015ApJ...813...93S", "2015MNRAS.446..683D", "2015MNRAS.450..905G", "2015Sci...347.1123K", "2016ApJ...819L...8K", "2016ApJ...822...78G" ]
[ "10.1051/0004-6361/201628925", "10.48550/arXiv.1607.01617" ]
1607
1607.01617_arXiv.txt
Supernovae (\sne) have proved to be exceptionally useful for different astrophysical and cosmological applications. For example, Type Ia supernovae (\sneia) have been used as distance indicators to show that the expansion rate of the Universe is accelerating \citep[see \eg][for a review]{2011ARNPS..61..251G}, while the rate of core-collapse \sne (CC~\sn) can be used to trace the star formation history (SFH; see \eg \citealt[]{2008ApJ...681..462D, 2012ApJ...757...70D}). Further, \sne contribute to the heavy elements in the Universe, so understanding the redshift dependence of \sn rates informs us about the chemical enrichment of galaxies over cosmic time. Measurements of SN rates at very high redshifts, $z \gtrsim 1$, are particularly difficult because of the limited light-gathering power of existing telescopes. This has been especially problematic for the study of CC~\sn rates, since they are on average intrinsically fainter than \sneia and often embedded in dusty environments \citep[see \eg][]{Mattila12}. Using ground-based facilities, only a few surveys have been able to measure CC~\sn rates beyond $z\gtrsim0.4$ \citep{Graur11,Melinder}. An advance in redshift was provided by the Hubble Space Telescope ({\it HST}) by extending the CC~SN rates measurements up to $z\approx2.5$ \citep{2004ApJ...613..189D,2012ApJ...757...70D,Strolger15} . This progress, however, was only made possible by the advent of very large {\it HST} multi-cycle treasury programmes: the Great Observatories Origins Deep Survey (GOODS), the Cosmic Assembly Near-Infrared Deep Extragalactic Legacy Survey (CANDELS) and the Cluster Lensing And Supernova survey with Hubble (CLASH). The magnification power of galaxy clusters can be used as gravitational telescopes to enhance survey depth \citep{1988ApJ...332...75N,1988ApJ...335L...9K,1998MNRAS.296..763K, 2000ApJ...532..679P,2003A&A...405..859G}. Galaxy clusters are the most massive gravitationally bound objects in the Universe, distorting and magnifying objects behind them. Gravitational lensing magnifies both the area and the flux of background objects, thereby increasing the depth of the survey; also the ability to find very high-$z$ \sne is enhanced. However, conservation of flux ensures that the source-plane area behind the cluster is shrunk owing to lensing, implying that the expected number of \sne does not necessarily increase in the field of view \citep[see][for a description of possible optimizations]{2003A&A...405..859G}. Even though \citet{1937ApJ....86..217Z} suggested the use of gravitational telescopes nearly 80 years ago, it is only recently that systematic \sn searches have been performed in background galaxies behind clusters. These investigations were initiated by feasibility studies by \citet{2000MNRAS.319..549S} and \citet{2003A&A...405..859G}. \citet{2002MNRAS.332...37G} searched {\it HST} archival images of galaxy cluster fields for lensed \sne, finding one \sn candidate at $z=0.985$. With its 524-orbit survey aimed at 25 galaxy clusters, one of the main objectives of the CLASH survey was to find lensed \sne behind the clusters \citep{2012ApJS..199...25P}. Three transients, out of which two were classified as secure \sneia, were detected and used as direct tests of independently derived lensing magnification maps \citep{2014MNRAS.440.2742N,2014ApJ...786....9P}. The Frontier Fields survey\footnote{www.stsci.edu/hst/campaigns/frontier-fields/} is an ongoing effort devoting almost a thousand {\it HST} orbits and targeting six lensing galaxy clusters. One of its most remarkable discoveries was a \snia at $z=1.346$ behind Abell 2744 \citep{2015arXiv150506211R}. Strong lensing also gives multiple images of the galaxies behind the cluster that host \sne. Even though the probability of observing such events is very low, multiple images of a strongly lensed \sn from Grism Lens-Amplified Survey from Space (GLASS; \citealt{2015ApJ...812..114T}) were detected. GLASS is a complementary {\it HST} spectroscopic survey targeting ten clusters, including those covered by Frontier Fields. The \sn was at $z=1.489$ (dubbed 'SN Refsdal') behind MACS J1149.6+2223, which re-appeared almost a year later \citep{2015ApJ...811...70R,2015Sci...347.1123K, 2015arXiv151104093G, 2015arXiv151204654K}. The ground-based near-infrared (NIR) search for lensed \sne behind galaxy clusters was pioneered by \citet{First} and \citet{Second} (hereafter G09). A search using the ISAAC instrument at the Very Large Telescope (VLT) was carried out targeting Abell~1689, Abell~1835, and AC114, with observations separated by one month. For Abell~1689, the gravitational telescope used in this work, a total of six hours of observations in $SZ$ band (similar to $Y$ band) were used, along with archival data used as a reference for image subtractions. This survey resulted in the discovery of one reddened, highly magnified SN~IIP at $z=0.6$ with a high lensing magnification from archival data taken in 2003. In addition to searching for transients in lensed galaxies, monitoring galaxy clusters offers the opportunity to detect \sne that originate from cluster members. These are mostly thermonuclear \sne, since clusters are dominated by early-type galaxies. Cluster \snia rates have been proposed to help disentangle the proposed scenarios for \snia progenitors and are essential in understanding several astrophysical processes such as the iron abundance in intracluster medium \citep{2007ApJ...660.1165S,2008AJ....135.1343G,2008MNRAS.383.1121M, 2010ApJ...722.1879M, 2010ApJ...713.1026D,2012ApJ...745...32B}. Here, we present the continuation of the effort of \citet{First} and \citet{Second} with HAWK-I on the VLT, which has greater sensitivity and a wider field of view than ISAAC. This paper is organized as follows. In Sect.\,\ref{sec:surveys} our surveys are presented and in Sect.\,3 the transient search strategy and \sn candidates are described. In Sect.\,4 the volumetric \sn rates are calculated; Sect.\,5 regards the connection between CC~\sn rates and SFH. Sect.\,6 concerns the rates in the galaxies with multiple images, while in Sect.\,7 the cluster \snia rates are calculated. In Sect.\,8 the possibilities of future surveys detecting \sne in the strongly lensed galaxies behind galaxy clusters are discussed. Also, given the lack of \sne at $z\gtrsim 2$, the requirements of finding very high-$z$ \sne are investigated. In Sect.\,9, summary and conclusions are drawn. Throughout the paper we assume the cosmology $\Omega_{\Lambda}=0.7$, $\Omega_{\rm M}=0.3$ and $h=0.7$, unless stated otherwise. All the magnitudes are given in the Vega system.
In this work we present the results of a dedicated ground-based NIR rolling search, accompanied with an optical programme, which is aimed at discovering high-$z$ lensed \sne behind the galaxy cluster Abell~1689. During 2008--2014, we obtained a total of 29~and 19~epochs in the \Jband and $i$ bands, respectively, and discovered five~CC~\sne behind the cluster and two\sneia associated with A1689. Notably, one of the most distant CC~\sn ever discovered was found at $z=1.703$ with significant magnification. Using these discoveries, we compute the volumetric CC~\sn rates in three redshift bins in the range $0.4<z<1.9$, and put upper limits on the CC~\sn rates in two additional redshift bins in $1.9<z<2.9$. Upper limits of the volumetric \snia rate are also calculated for the same redshift bins. All the measured rates are found to be consistent with previous studies at the corresponding redshifts. We further emphasize the comparably high statistical precision at high redshift, given the modest investment in observing time, which can be obtained using gravitational telescopes. The impact of systematic uncertainties were calculated for the CC~\sn rates, which will become increasingly important for upcoming wide-field surveys such as the LSST which are expected to discover a large number of \sne \citep{2009JCAP...01..047L}. We highlight the need for better understanding CC~SN properties, such as the subtype fractions and peak magnitudes, which will affect the systematics budget. Since the CC~SN rate traces the cosmic SFH, we compare our results and literature values to the expected rates calculated from the \citet{2014ARA&A..52..415M} SFH. We find that the measured and predicted rates are in good agreement with a scale factor of, $k_{\rm CC}=0.0093\pm0.0010$~$\msun^{-1}$. By simulating ground- and space-based five-year surveys, we explore the possibility of finding lensed \sne at $z\gtrsim 2.0$. We find that there is very little room for improvement with the current facilities and that the next generation of telescopes, for example WFIRST, are needed. Finally, we estimate the number of strongly lensed \sne with multiple images that can be expected to be discovered behind A1689 by upcoming transient surveys. We find that LSST, and in particular WFIRST, can be expected to find tens of strongly lensed SNe that would allow the time delays between the multiple images to be measured. Until the first light of these surveys, using gravitational telescopes is the only way to find high-$z$ \sne using available ground-based instruments.
16
7
1607.01617
<BR /> Aims: We present a ground-based, near-infrared search for lensed supernovae behind the massive cluster Abell 1689 at z = 0.18, which is one of the most powerful gravitational telescopes that nature provides. <BR /> Methods: Our survey was based on multi-epoch J-band observations with the HAWK-I instrument on VLT, with supporting optical data from the Nordic Optical Telescope. <BR /> Results: Our search resulted in the discovery of five photometrically classified, core-collapse supernovae with high redshifts of 0.671 &lt; z &lt; 1.703 and magnifications in the range Δm = - 0.31 to -1.58 mag, as calculated from lensing models in the literature. Owing to the power of the lensing cluster, the survey had the sensitivity to detect supernovae up to very high redshifts, z 3, albeit for a limited region of space. We present a study of the core-collapse supernova rates for 0.4 ≤ z&lt; 2.9, and find good agreement with previous estimates and predictions from star formation history. During our survey, we also discovered two Type Ia supernovae in A 1689 cluster members, which allowed us to determine the cluster Ia rate to be 0.14<SUP>+0.19</SUP><SUB>-0.09</SUB>±0.01SNuB h<SUP>2</SUP> (SNuB≡10<SUP>-12</SUP>SNe L<SUP>-1</SUP><SUB>⊙,B</SUB> yr<SUP>-1</SUP>), where the error bars indicate 1σ confidence intervals, statistical and systematic, respectively. The cluster rate normalized by the stellar mass is 0.10<SUP>+0.13</SUP><SUB>-0.096</SUB>±0.02 in SNuM h<SUP>2</SUP> (SNuM ≡10<SUP>-12</SUP>SNe M<SUP>-1</SUP><SUB>⊙</SUB> yr<SUP>-1</SUP>). Furthermore, we explore the optimal future survey for improving the core-collapse supernova rate measurements at z ≳ 2 using gravitational telescopes, and for detections with multiply lensed images, and we find that the planned WFIRST space mission has excellent prospects. <BR /> Conclusions: Massive clusters can be used as gravitational telescopes to significantly expand the survey range of supernova searches, with important implications for the study of the high-z transient Universe. <P />Based on observations made with European Southern Observatory (ESO) telescopes at the Paranal Observatory under programme ID 082.A-0431; 0.83.A-0398, 091.A-0108 and ID 093.A-0278, PI: A. Goobar.The deep average image (FITS file) is 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/594/A54">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/594/A54</A>
false
[ "excellent prospects", "star formation history", "A. Goobar", "gravitational telescopes", "high redshifts", "1σ confidence intervals", "multiply lensed images", "programme ID", "lensing models", "Massive clusters", "anonymous ftp", "supernova searches", "European Southern Observatory", "space", "lensed supernovae", "SNuM h", "optical data", "z", "previous estimates", "good agreement" ]
2.725525
4.552824
-1
5398349
[ "Aartsen, M. G.", "Abraham, K.", "Ackermann, M.", "Adams, J.", "Aguilar, J. A.", "Ahlers, M.", "Ahrens, M.", "Altmann, D.", "Andeen, K.", "Anderson, T.", "Ansseau, I.", "Anton, G.", "Archinger, M.", "Argüelles, C.", "Auffenberg, J.", "Axani, S.", "Bai, X.", "Barwick, S. W.", "Baum, V.", "Bay, R.", "Beatty, J. J.", "Becker Tjus, J.", "Becker, K. -H.", "BenZvi, S.", "Berghaus, P.", "Berley, D.", "Bernardini, E.", "Bernhard, A.", "Besson, D. Z.", "Binder, G.", "Bindig, D.", "Bissok, M.", "Blaufuss, E.", "Blot, S.", "Bohm, C.", "Börner, M.", "Bos, F.", "Bose, D.", "Böser, S.", "Botner, O.", "Braun, J.", "Brayeur, L.", "Bretz, H. -P.", "Burgman, A.", "Carver, T.", "Casier, M.", "Cheung, E.", "Chirkin, D.", "Christov, A.", "Clark, K.", "Classen, L.", "Coenders, S.", "Collin, G. H.", "Conrad, J. M.", "Cowen, D. F.", "Cross, R.", "Day, M.", "de André, J. P. A. M.", "De Clercq, C.", "del Pino Rosendo, E.", "Dembinski, H.", "De Ridder, S.", "Desiati, P.", "de Vries, K. D.", "de Wasseige, G.", "de With, M.", "DeYoung, T.", "Díaz-Vélez, J. C.", "di Lorenzo, V.", "Dujmovic, H.", "Dumm, J. P.", "Dunkman, M.", "Eberhardt, B.", "Ehrhardt, T.", "Eichmann, B.", "Eller, P.", "Euler, S.", "Evenson, P. A.", "Fahey, S.", "Fazely, A. R.", "Feintzeig, J.", "Felde, J.", "Filimonov, K.", "Finley, C.", "Flis, S.", "Fösig, C. -C.", "Franckowiak, A.", "Friedman, E.", "Fuchs, T.", "Gaisser, T. K.", "Gallagher, J.", "Gerhardt, L.", "Ghorbani, K.", "Giang, W.", "Gladstone, L.", "Glagla, M.", "Glüsenkamp, T.", "Goldschmidt, A.", "Golup, G.", "Gonzalez, J. G.", "Grant, D.", "Griffith, Z.", "Haack, C.", "Haj Ismail, A.", "Hallgren, A.", "Halzen, F.", "Hansen, E.", "Hansmann, B.", "Hansmann, T.", "Hanson, K.", "Hebecker, D.", "Heereman, D.", "Helbing, K.", "Hellauer, R.", "Hickford, S.", "Hignight, J.", "Hill, G. C.", "Hoffman, K. D.", "Hoffmann, R.", "Holzapfel, K.", "Hoshina, K.", "Huang, F.", "Huber, M.", "Hultqvist, K.", "In, S.", "Ishihara, A.", "Jacobi, E.", "Japaridze, G. S.", "Jeong, M.", "Jero, K.", "Jones, B. J. P.", "Jurkovic, M.", "Kappes, A.", "Karg, T.", "Karle, A.", "Katz, U.", "Kauer, M.", "Keivani, A.", "Kelley, J. L.", "Kemp, J.", "Kheirandish, A.", "Kim, M.", "Kintscher, T.", "Kiryluk, J.", "Kittler, T.", "Klein, S. R.", "Kohnen, G.", "Koirala, R.", "Kolanoski, H.", "Konietz, R.", "Köpke, L.", "Kopper, C.", "Kopper, S.", "Koskinen, D. J.", "Kowalski, M.", "Krings, K.", "Kroll, M.", "Krückl, G.", "Krüger, C.", "Kunnen, J.", "Kunwar, S.", "Kurahashi, N.", "Kuwabara, T.", "Labare, M.", "Lanfranchi, J. L.", "Larson, M. J.", "Lauber, F.", "Lennarz, D.", "Lesiak-Bzdak, M.", "Leuermann, M.", "Leuner, J.", "Lu, L.", "Lünemann, J.", "Madsen, J.", "Maggi, G.", "Mahn, K. B. M.", "Mancina, S.", "Mandelartz, M.", "Maruyama, R.", "Mase, K.", "Maunu, R.", "McNally, F.", "Meagher, K.", "Medici, M.", "Meier, M.", "Meli, A.", "Menne, T.", "Merino, G.", "Meures, T.", "Miarecki, S.", "Mohrmann, L.", "Montaruli, T.", "Moulai, M.", "Nahnhauer, R.", "Naumann, U.", "Neer, G.", "Niederhausen, H.", "Nowicki, S. C.", "Nygren, D. R.", "Obertacke Pollmann, A.", "Olivas, A.", "O'Murchadha, A.", "Palczewski, T.", "Pandya, H.", "Pankova, D. V.", "Penek, Ö.", "Pepper, J. A.", "Pérez de los Heros, C.", "Pieloth, D.", "Pinat, E.", "Price, P. B.", "Przybylski, G. T.", "Quinnan, M.", "Raab, C.", "Rädel, L.", "Rameez, M.", "Rawlins, K.", "Reimann, R.", "Relethford, B.", "Relich, M.", "Resconi, E.", "Rhode, W.", "Richman, M.", "Riedel, B.", "Robertson, S.", "Rongen, M.", "Rott, C.", "Ruhe, T.", "Ryckbosch, D.", "Rysewyk, D.", "Sabbatini, L.", "Sanchez Herrera, S. E.", "Sandrock, A.", "Sandroos, J.", "Sarkar, S.", "Satalecka, K.", "Schimp, M.", "Schlunder, P.", "Schmidt, T.", "Schoenen, S.", "Schöneberg, S.", "Schumacher, L.", "Seckel, D.", "Seunarine, S.", "Soldin, D.", "Song, M.", "Spiczak, G. M.", "Spiering, C.", "Stahlberg, M.", "Stanev, T.", "Stasik, A.", "Steuer, A.", "Stezelberger, T.", "Stokstad, R. G.", "Stössl, A.", "Ström, R.", "Strotjohann, N. L.", "Sullivan, G. W.", "Sutherland, M.", "Taavola, H.", "Taboada, I.", "Tatar, J.", "Tenholt, F.", "Ter-Antonyan, S.", "Terliuk, A.", "Tešić, G.", "Tilav, S.", "Toale, P. A.", "Tobin, M. N.", "Toscano, S.", "Tosi, D.", "Tselengidou, M.", "Turcati, A.", "Unger, E.", "Usner, M.", "Vandenbroucke, J.", "van Eijndhoven, N.", "Vanheule, S.", "van Rossem, M.", "van Santen, J.", "Veenkamp, J.", "Vehring, M.", "Voge, M.", "Vraeghe, M.", "Walck, C.", "Wallace, A.", "Wallraff, M.", "Wandkowsky, N.", "Weaver, Ch.", "Weiss, M. J.", "Wendt, C.", "Westerhoff, S.", "Whelan, B. J.", "Wickmann, S.", "Wiebe, K.", "Wiebusch, C. H.", "Wille, L.", "Williams, D. R.", "Wills, L.", "Wolf, M.", "Wood, T. R.", "Woolsey, E.", "Woschnagg, K.", "Xu, D. L.", "Xu, X. W.", "Xu, Y.", "Yanez, J. P.", "Yodh, G.", "Yoshida, S.", "Zoll, M.", "IceCube Collaboration" ]
2016ApJ...830..129A
[ "Search for Sources of High-energy Neutrons with Four Years of Data from the IceTop Detector" ]
2
[ "Department of Physics, University of Adelaide, Adelaide, 5005, Australia", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "DESY, D-15735 Zeuthen, Germany", "Department of Physics and Astronomy, University of Canterbury, Private Bag 4800, Christchurch, New Zealand", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Erlangen Centre for Astroparticle Physics, Friedrich-Alexander-Universität Erlangen-Nürnberg, D-91058 Erlangen, Germany", "Department of Physics, Marquette University, Milwaukee, WI 53201, USA", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "Erlangen Centre for Astroparticle Physics, Friedrich-Alexander-Universität Erlangen-Nürnberg, D-91058 Erlangen, Germany", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Department of Physics, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Physics Department, South Dakota School of Mines and Technology, Rapid City, SD 57701, USA", "Department of Physics and Astronomy, University of California, Irvine, CA 92697, USA", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Department of Physics, University of California, Berkeley, CA 94720, USA", "Department of Physics and Center for Cosmology and Astro-Particle Physics, Ohio State University, Columbus, OH 43210, USA ; Department of Astronomy, Ohio State University, Columbus, OH 43210, USA", "Fakultät für Physik &amp; Astronomie, Ruhr-Universität Bochum, D-44780 Bochum, Germany", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "Department of Physics and Astronomy, University of Rochester, Rochester, NY 14627, USA", "National Research Nuclear University MEPhI (Moscow Engineering Physics Institute), Moscow, Russia", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "DESY, D-15735 Zeuthen, Germany", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "Department of Physics and Astronomy, University of Kansas, Lawrence, KS 66045, USA", "Department of Physics, University of California, Berkeley, CA 94720, USA; Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "DESY, D-15735 Zeuthen, Germany", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Fakultät für Physik &amp; Astronomie, Ruhr-Universität Bochum, D-44780 Bochum, Germany", "Department of Physics, Sungkyunkwan University, Suwon 440-746, Korea", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "DESY, D-15735 Zeuthen, Germany", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "Département de physique nucléaire et corpusculaire, Université de Genève, CH-1211 Genève, Switzerland", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Département de physique nucléaire et corpusculaire, Université de Genève, CH-1211 Genève, Switzerland", "Department of Physics, University of Toronto, Toronto, ON M5S 1A7, Canada", "Institut für Kernphysik, Westfälische Wilhelms-Universität Münster, D-48149 Münster, Germany", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "Department of Physics, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Department of Physics, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA; Department of Astronomy and Astrophysics, Pennsylvania State University, University Park, PA 16802, USA", "Department of Physics and Astronomy, University of Rochester, Rochester, NY 14627, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Department of Physics and Astronomy, University of Gent, B-9000 Gent, Belgium", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Institut für Physik, Humboldt-Universität zu Berlin, D-12489 Berlin, Germany", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Department of Physics, Sungkyunkwan University, Suwon 440-746, Korea", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Fakultät für Physik &amp; Astronomie, Ruhr-Universität Bochum, D-44780 Bochum, Germany", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, Southern University, Baton Rouge, LA 70813, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics, University of California, Berkeley, CA 94720, USA", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "DESY, D-15735 Zeuthen, Germany", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Department of Astronomy, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, University of California, Berkeley, CA 94720, USA; Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "DESY, D-15735 Zeuthen, Germany", "Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics and Astronomy, University of Gent, B-9000 Gent, Belgium", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Niels Bohr Institute, University of Copenhagen, DK-2100 Copenhagen, Denmark", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Institut für Physik, Humboldt-Universität zu Berlin, D-12489 Berlin, Germany", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Department of Physics, University of Adelaide, Adelaide, 5005, Australia", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA; Earthquake Research Institute, University of Tokyo, Bunkyo, Tokyo 113-0032, Japan", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Department of Physics, Sungkyunkwan University, Suwon 440-746, Korea", "Department of Physics, Chiba University, Chiba 263-8522, Japan", "DESY, D-15735 Zeuthen, Germany", "CTSPS, Clark-Atlanta University, Atlanta, GA 30314, USA", "Department of Physics, Sungkyunkwan University, Suwon 440-746, Korea", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "Institut für Kernphysik, Westfälische Wilhelms-Universität Münster, D-48149 Münster, Germany", "DESY, D-15735 Zeuthen, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Erlangen Centre for Astroparticle Physics, Friedrich-Alexander-Universität Erlangen-Nürnberg, D-91058 Erlangen, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, Sungkyunkwan University, Suwon 440-746, Korea", "DESY, D-15735 Zeuthen, Germany", "Department of Physics and Astronomy, Stony Brook University, Stony Brook, NY 11794-3800, USA", "Erlangen Centre for Astroparticle Physics, Friedrich-Alexander-Universität Erlangen-Nürnberg, D-91058 Erlangen, Germany", "Department of Physics, University of California, Berkeley, CA 94720, USA; Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Université de Mons, B-7000 Mons, Belgium", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Institut für Physik, Humboldt-Universität zu Berlin, D-12489 Berlin, Germany", "III. 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Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics, Chiba University, Chiba 263-8522, Japan", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Department of Physics, University of Wisconsin, River Falls, WI 54022, USA", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Fakultät für Physik &amp; Astronomie, Ruhr-Universität Bochum, D-44780 Bochum, Germany", "Department of Physics, Yale University, New Haven, CT 06520, USA", "Department of Physics, Chiba University, Chiba 263-8522, Japan", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "Niels Bohr Institute, University of Copenhagen, DK-2100 Copenhagen, Denmark", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Department of Physics and Astronomy, University of Gent, B-9000 Gent, Belgium", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "Department of Physics, University of California, Berkeley, CA 94720, USA; Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "DESY, D-15735 Zeuthen, Germany", "Département de physique nucléaire et corpusculaire, Université de Genève, CH-1211 Genève, Switzerland", "Department of Physics, Massachusetts Institute of Technology, Cambridge, MA 02139, USA", "DESY, D-15735 Zeuthen, Germany", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Department of Physics and Astronomy, Stony Brook University, Stony Brook, NY 11794-3800, USA", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "Department of Physics and Astronomy, University of Alabama, Tuscaloosa, AL 35487, USA", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics and Astronomy, University of Alabama, Tuscaloosa, AL 35487, USA", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "Department of Physics, University of California, Berkeley, CA 94720, USA", "Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Université Libre de Bruxelles, Science Faculty CP230, B-1050 Brussels, Belgium", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Niels Bohr Institute, University of Copenhagen, DK-2100 Copenhagen, Denmark", "Department of Physics and Astronomy, University of Alaska Anchorage, 3211 Providence Drive, Anchorage, AK 99508, USA", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics, Drexel University, 3141 Chestnut Street, Philadelphia, PA 19104, USA", "Department of Physics, Chiba University, Chiba 263-8522, Japan", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Department of Physics, Drexel University, 3141 Chestnut Street, Philadelphia, PA 19104, USA", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Department of Physics, University of Adelaide, Adelaide, 5005, Australia", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics, Sungkyunkwan University, Suwon 440-746, Korea", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Department of Physics and Astronomy, University of Gent, B-9000 Gent, Belgium", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Niels Bohr Institute, University of Copenhagen, DK-2100 Copenhagen, Denmark; Department of Physics, University of Oxford, 1 Keble Road, Oxford OX1 3NP, UK", "DESY, D-15735 Zeuthen, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics, TU Dortmund University, D-44221 Dortmund, Germany", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Fakultät für Physik &amp; Astronomie, Ruhr-Universität Bochum, D-44780 Bochum, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Department of Physics, University of Wisconsin, River Falls, WI 54022, USA", "Department of Physics, University of Wuppertal, D-42119 Wuppertal, Germany", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics, University of Wisconsin, River Falls, WI 54022, USA", "DESY, D-15735 Zeuthen, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "DESY, D-15735 Zeuthen, Germany", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "DESY, D-15735 Zeuthen, Germany", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "DESY, D-15735 Zeuthen, Germany", "Department of Physics, University of Maryland, College Park, MD 20742, USA", "Department of Physics and Center for Cosmology and Astro-Particle Physics, Ohio State University, Columbus, OH 43210, USA", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "School of Physics and Center for Relativistic Astrophysics, Georgia Institute of Technology, Atlanta, GA 30332, USA", "Department of Physics, University of California, Berkeley, CA 94720, USA; Lawrence Berkeley National Laboratory, Berkeley, CA 94720, USA", "Fakultät für Physik &amp; Astronomie, Ruhr-Universität Bochum, D-44780 Bochum, Germany", "Department of Physics, Southern University, Baton Rouge, LA 70813, USA", "DESY, D-15735 Zeuthen, Germany", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Bartol Research Institute and Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA", "Department of Physics and Astronomy, University of Alabama, Tuscaloosa, AL 35487, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Erlangen Centre for Astroparticle Physics, Friedrich-Alexander-Universität Erlangen-Nürnberg, D-91058 Erlangen, Germany", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "Department of Physics and Astronomy, Uppsala University, Box 516, S-75120 Uppsala, Sweden", "DESY, D-15735 Zeuthen, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Vrije Universiteit Brussel, Dienst ELEM, B-1050 Brussels, Belgium", "Department of Physics and Astronomy, University of Gent, B-9000 Gent, Belgium", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "DESY, D-15735 Zeuthen, Germany", "Physik-department, Technische Universität München, D-85748 Garching, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Physikalisches Institut, Universität Bonn, Nussallee 12, D-53115 Bonn, Germany", "Department of Physics and Astronomy, University of Gent, B-9000 Gent, Belgium", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Department of Physics, University of Adelaide, Adelaide, 5005, Australia", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Department of Physics, Pennsylvania State University, University Park, PA 16802, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, University of Adelaide, Adelaide, 5005, Australia", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Institute of Physics, University of Mainz, Staudinger Weg 7, D-55099 Mainz, Germany", "III. Physikalisches Institut, RWTH Aachen University, D-52056 Aachen, Germany", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics and Astronomy, University of Alabama, Tuscaloosa, AL 35487, USA", "Department of Physics, Drexel University, 3141 Chestnut Street, Philadelphia, PA 19104, USA", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Department of Physics, University of Alberta, Edmonton, AB T6G 2E1, Canada", "Department of Physics, University of California, Berkeley, CA 94720, USA", "Department of Physics and Wisconsin IceCube Particle Astrophysics Center, University of Wisconsin, Madison, WI 53706, USA", "Department of Physics, Southern University, Baton Rouge, LA 70813, USA", "Department of Physics and Astronomy, Stony Brook University, Stony Brook, NY 11794-3800, USA", "DESY, D-15735 Zeuthen, Germany", "Department of Physics and Astronomy, University of California, Irvine, CA 92697, USA", "Department of Physics, Chiba University, Chiba 263-8522, Japan", "Oskar Klein Centre and Department of Physics, Stockholm University, SE-10691 Stockholm, Sweden", "-" ]
[ "2017ApJ...837L..25A", "2018AdSpR..62.2902A" ]
[ "astronomy", "physics" ]
6
[ "astroparticle physics", "cosmic rays", "methods: data analysis", "Astrophysics - High Energy Astrophysical Phenomena", "High Energy Physics - Experiment" ]
[ "1983ApJ...272..317L", "1989A&A...213L..23S", "1997PhRvL..79.1805C", "1997PhRvL..79.2616B", "1998ApJ...493..175B", "1998PhRvD..57.3873F", "1999APh....10..303H", "2001APh....15..167B", "2002APh....17...23C", "2003JPhG...29.1409B", "2004ApJ...608..865A", "2005AJ....129.1993M", "2005ApJ...622..759G", "2005ApJ...622..892C", "2005IJMPA..20.6753H", "2006APh....26...41C", "2007A&A...469..807L", "2009ARA&A..47..523H", "2009NIMPA.601..294A", "2010NIMPA.618..139A", "2011A&A...527L...4B", "2011arXiv1103.4284N", "2012ApJ...760..148P", "2012MNRAS.424.2249K", "2013APh....47...54A", "2013ApJ...765...55A", "2013ApJS..208...17A", "2013NIMPA.700..188A", "2013PhRvD..87f2002A", "2013PhRvD..88d2004A", "2014ApJ...789L..34A", "2014CRPhy..15..329B", "2014NuPhS.256..149B", "2015ApJ...804...15A", "2015ApJ...804..133A" ]
[ "10.3847/0004-637X/830/2/129", "10.48550/arXiv.1607.05614" ]
1607
1607.05614_arXiv.txt
\label{sec:intro} The Galactic magnetic field (GMF) strongly affects the arrival distribution of charged cosmic rays, thereby obscuring their sources. A compact source of high energy neutrons would manifest as a point source in cosmic ray arrival directions since neutrons are not deflected by magnetic fields. Secondary neutral particles are an expected signature of hadronic acceleration in Galactic sources. Neutral particles would be produced as the cosmic ray protons and nuclei undergo $pp$ and $p\gamma$ collisions, and photodisintegration, respectively, on the ambient photons and cosmic rays within the dense environment surrounding their source (see, e.g., \citep{2002APh....17...23C, Crocker:2005bb, 2006APh....26...41C, 2007PhRvD..75f3001A}). For example, neutrons result from charge-exchange interactions, \begin{equation*} p\gamma \rightarrow n \pi^{+} \end{equation*} where a $\pi^{+}$ emerges with the proton's positive charge and the neutron retains most of the energy. For interacting proton primaries, photons resulting from $\pi^{0}$ decays take a small fraction of the proton energy. The production of neutrons exceeds the production of photons at the same energy \citep{Crocker:2005bb}.% It is plausible that known Galactic sources could produce high energy neutron fluxes, based on the measured TeV energy photon flux. For some Galactic sources, the energy flux of TeV photons is greater than 1 eV cm$^{-2}$ s$^{-1}$ \citep{hinton2009}. Sources producing particle fluxes with an $E^{-2}$ differential energy spectrum inject equal energy into each energy decade. If sources in the Galaxy produce PeV photons in addition to TeV photons, the PeV photon energy flux would also exceed 1 eV cm$^{-2}$ s$^{-1}$ at Earth. For sources that produce neutrons by hadronic processes as well, the neutron energy flux would be even higher since the neutron production rate exceeds the photon production rate, as noted previously. Free neutrons undergo beta decay with a $880.0\pm0.9$ second half-life \citep{PDG:2014}. Due to this decay, sources will only be visible within about 10 ($E$ / PeV) pc of Earth. Since plausible accelerators such as young pulsars are no closer than 100 pc, searches at energies above 10 PeV are the most promising. A diffuse flux of neutrons could be expected from interactions of cosmic ray primaries with ambient photons and the interstellar medium. However, at PeV energies this flux would appear all over the sky since the effective range is less than the thickness of the Galactic disk. This complicates a search for correlations with the Galactic plane since an excess signal could not be constrained to a particular region of the sky, for example Galactic latitudes $\abs{b}<10^{\circ}$. At energies above $10^{18}$ eV (1 EeV), the Pierre Auger Observatory recently performed a search for neutrons in the Southern hemisphere finding no significant signal excesses or correlations with catalogs of Galactic objects, and established flux upper limits \citep{Aab:2012bha, Aab:2014caa}. The Telescope Array experiment has established flux limits for point sources above 0.5 EeV in the Northern hemisphere \citep{2015ApJ...804..133A}. KASCADE \citep{Antoni:2004sc} and CASA-MIA \citep{Chantell1997, Borione1998} found no point sources in the Northern hemisphere, also setting flux limits (an all-sky limit in the case of KASCADE). AGASA \citep{Hayashida1999} and a re-analysis \citep{Bellido2001} of SUGAR data reported slight excesses towards the Galactic center, although these were later not confirmed by Auger \citep{Aab:2014yba}. This paper reports the results of two searches for point-like signals in the arrival direction distribution of four years of IceTop data. The two searches are an all-sky search for general hotspots on the sky and a search for correlations with nearby known Galactic sources. In the all-sky search, we look for an excess of events from any direction in the sky, evaluating the significance of any excess using the method of Li and Ma \citep{Li:1983fv}. The observable signature of a neutron flux is an excess of proton-like air showers. The targeted search is treated as a stacked analysis using a set of candidate sources from an astrophysical catalog. It is assumed that many or all of the candidates for a given set are emitting neutrons, so the combined signal should be more significant than that of a single target. In both the all-sky and targeted searches, we set flux upper limits using the procedures of Feldman and Cousins \citep{Feldman:1997qc}. This paper is organized as follows. In Section \ref{sec:IC-IT}, the IceTop detector is described. Section \ref{sec:recmethdata} summarizes the reconstruction methods and characteristics of the dataset. The analysis methods and details of the search methods are described in Section \ref{sec:searches}. The search results are presented in Section \ref{sec:results}. A discussion of the results (Section \ref{sec:summary}) concludes the paper.
\label{sec:summary} IceTop does not observe a statistically significant point source of cosmic ray arrival directions. Using Equation \ref{eq:conversionfactor} the all-sky mean flux upper limits for individual declination bands correspond to energy fluxes between about 0.6 - 1.2 eV cm$^{-2}$ s$^{-1}$ between 100 PeV and 1 EeV assuming an $E^{-2}$ neutron energy spectrum as measured at Earth, which are comparable to TeV photon fluxes for Galactic objects \citep{hinton2009}. These flux limits are the first neutron flux upper limits in the Southern hemisphere for energies in the 10 PeV to 1 EeV energy decades. Again, it is important to note that neutron decay en-route will modify the energy spectrum as illustrated in Figure \ref{fig:decayatten}, so the source spectrum would be generally softer than that constrained. The limits in both searches are strongly dependent on the assumption that an injected $E^{-2}$ spectrum is not significantly modified by decay, as noted in Section \ref{sec:fulcalc}. For the all-sky search, this restricts the applicability of the limits within a small volume around Earth. For the targeted search, there are a number of objects that lie within 1 kpc, so their limits are most compatible with the base assumption. As noted previously, hadronic production of photons by protons with an $E^{-2}$ spectrum will inject equal power into each energy decade, and the neutron production at least equals the photon production. At present, these flux upper limits do not strongly constrain the TeV photon production mechanism, or the shape of the parent energy spectrum. No significant correlation is found with known nearby Galactic objects characterized by GeV-TeV energy photon emission and plausibly capable of producing PeV neutrons. The non-observation of PeV neutrons may simply indicate that these objects are not producing neutrons at these energies, or that typical Galactic neutron sources are not near Earth. Local PeV neutron production in the Galaxy could simply be episodic or transient, for example, occurring during supernova explosions or other extremely high energy particle production events. Alternatively, the sources may emit particle jets continuously, but their number may be few and the jets are not oriented towards the Earth. Individual sources could emit weakly but be densely distributed. Additionally, the environment around any sources may not be sufficiently dense to facilitate neutron production by cosmic ray interaction such that the primaries escape the acceleration region into interstellar space before interacting and producing neutrons. In this case, neutrons decay in interstellar space relatively near the primary source producing secondary protons \citep{1997PhRvL..79.2616B}. These secondary protons then propagate diffusively in the GMF, so sources that are sufficiently far away will not manifest a point source signal of cosmic ray neutrons, but could contribute to a proton signal that is smeared on the sky and not necessarily pointing back to the original source; this argument was presented by \citep{Bossa2003} when they considered EeV neutrons from the Galactic center. At PeV energies, neutrons would penetrate much less into the surrounding medium, so any potential signal from the resulting protons would be strongly suppressed by the scattering effects of the GMF and masked by the background cosmic ray flux. At higher energies, for example, between 10-100 PeV, this process could further enrich the cosmic ray proton fraction above that which is directly accelerated at the source. The knee in the cosmic ray spectrum is observed around 4 PeV which is interpreted as an indication of a maximum attainable rigidity of typical Galactic cosmic ray sources and of associated changes in elemental composition (see e.g., \citep{Horandel:2005bb, Blasi:2014roa}). It is plausible that the maximum attainable energy for the proton energy spectrum in nearby sources may not extend well above the knee energy although for heavier compositions this scales with the nuclear charge $Z$. Above 10 PeV, the cosmic ray flux becomes progressively heavier with energy and with a decreasing proton fraction which is roughly 20\% at 10 PeV \citep{IceTopICRC2015-795, Apel:2013dga}. This suggests that such secondary enrichment may be unlikely since a recovering proton fraction is not observed at energies between roughly 10 PeV to a few 100 PeV. The non-observation of a PeV neutron flux does not necessarily preclude the existence of a PeV photon flux. The neutron energy spectrum at lower energies becomes increasingly modified by decay. PeV photons, on the other hand, have an absorption length considerably larger than the neutron decay distance and will maintain an unmodified energy spectrum that more resembles the injected spectrum at source. PeV photons could still plausibly be produced by non-hadronic processes, such as inverse-Compton scattering from a high energy electron population in or near Galactic sources (see e.g., \citep{1989AA...213L..23S, 2011arXiv1103.4284N, 2011A&A...527L...4B, 2012MNRAS.424.2249K}), although there are flux upper limits in the Northern \citep{Chantell1997, Borione1998, Feng_KG_2015, Kang_KG_2015a, Kang_KG_2015b} and Southern \citep{2013PhRvD..87f2002A} hemispheres. These photon limits, except for \citep{Kang_KG_2015a}, are for energies of order 1 PeV or below, whereas this analysis is most sensitive at energies above 100 PeV.
16
7
1607.05614
IceTop is an air-shower array located on the Antarctic ice sheet at the geographic South Pole. IceTop can detect an astrophysical flux of neutrons from Galactic sources as an excess of cosmic-ray air showers arriving from the source direction. Neutrons are undeflected by the Galactic magnetic field and can typically travel 10 (E/PeV) pc before decay. Two searches are performed using 4 yr of the IceTop data set to look for a statistically significant excess of events with energies above 10 PeV (10<SUP>16</SUP> eV) arriving within a small solid angle. The all-sky search method covers from -90° to approximately -50° in declination. No significant excess is found. A targeted search is also performed, looking for significant correlation with candidate sources in different target sets. This search uses a higher-energy cut (100 PeV) since most target objects lie beyond 1 kpc. The target sets include pulsars with confirmed TeV energy photon fluxes and high-mass X-ray binaries. No significant correlation is found for any target set. Flux upper limits are determined for both searches, which can constrain Galactic neutron sources and production scenarios.
false
[ "confirmed TeV energy photon fluxes", "different target sets", "Galactic neutron sources", "South Pole", "Galactic sources", "most target objects", "candidate sources", "energies", "PeV", "Flux upper limits", "high-mass X-ray binaries", "production scenarios", "decay", "significant correlation", "TeV", "Galactic", "pulsars", "cosmic-ray air showers", "declination", "the geographic South Pole" ]
5.700264
-0.50109
13
12409088
[ "Kisaka, Shota", "Asano, Katsuaki", "Terasawa, Toshio" ]
2016ApJ...829...12K
[ "Electric Field Screening with Backflow at Pulsar Polar Cap" ]
5
[ "Department of Physics and Mathematics, Aoyama Gakuin University, Sagamihara, Kanagawa, 252-5258, Japan ;;", "Institute for Cosmic Ray Research, The University of Tokyo, Kashiwa, Chiba, 277-8582, Japan", "Institute for Cosmic Ray Research, The University of Tokyo, Kashiwa, Chiba, 277-8582, Japan; iTHES Research Group, RIKEN, Wako, Saitama 351-0198, Japan;" ]
[ "2016ApJ...832..212M", "2017ApJ...837...76K", "2018MNRAS.475.5585Z", "2019MNRAS.483.4175Y", "2020ApJ...902...80K" ]
[ "astronomy" ]
8
[ "acceleration of particles", "pulsars: general", "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.3847/0004-637X/829/1/12", "10.48550/arXiv.1607.02202" ]
1607
1607.02202_arXiv.txt
\label{intro} In pulsar magnetosphere, particles are significantly accelerated at the given regions, and emit electromagnetic radiation from radio to $\gamma$-ray wavelength. Observations by {\it Fermi Gamma-Ray Space Telescope} have shown that the differential spectra above 200 MeV are well described by the power-law functions with an exponential cut-off, and that the cutoff shapes sharper than the simple exponential cutoff are rejected with high significance \citep[e.g., ][]{vela09}. This rules out the near-surface emission proposed in the polar cap cascade model \citep{DH96}, which would exhibit a much sharper spectral cut-off due to the attenuation of the magnetic pair-creation. Hence, the detected $\gamma$-ray pulse emission should originate from the outer region of the magnetosphere, as considered in the outer gap model \citep[e.g., ][]{CHR86, R96, THSC06, H06, H15, TNC15}, as well as the current sheet model \citep[e.g., ][]{KSG02, BS10, KHK14, BKHK15, CPS15}. On the other hand, the region just above the neutron star surface (NSS) has been considered as the site of the radio pulsed emission \citep[e.g., ][]{Nou+15}. The mechanism of pulsar radio emission is established as a coherent process, so that the plasma dynamics near the NSS would be strongly related to the emission mechanism \citep[e.g., ][]{S71}. The two-stream instability is a promising process to create the plasma bunches \citep[e.g., ][]{RS75}. The curvature radiation from the bunches is usually discussed as the mechanism of the coherent radio emission \citep[e.g., ][]{S75}. In order to investigate the possible instabilities near the NSS, we should take into account the non-stationary effects in the plasma flows. The dynamics of the plasma and the electromagnetic field near the NSS highly depends on the ratio of the current density parameter along the magnetic field, $j_{\rm m}$, to the Goldreich-Julian (GJ) value, $j_{\rm GJ}=\rho_{\rm GJ}c$ \citep{Me85, S97, TA13}, which is characterized by the GJ charge density $\rho_{\rm GJ}=-{\bf \Omega}\cdot{\bf B}/2\pi c$ \citep{GJ69}, where ${\bf \Omega}$ is the stellar angular velocity vector and ${\bf B}$ is the local magnetic field vector. The parameter $j_{\rm m}$ is regulated by the twist of the magnetic field ($\nabla\times{\bf B}$) imposed by the global stress balance of the pulsar magnetosphere \citep[e.g., ][]{S91}. In the polar cap region near the NSS, an accelerating electric field spontaneously develops to adjust the current and charge densities to the current density parameter $j_{\rm m}$ and the GJ charge density $\rho_{\rm GJ}$. In the cases $j_{\rm m}/j_{\rm GJ}\le0$ and $j_{\rm m}/j_{\rm GJ}\ge1$, outflowing particles from the NSS alone cannot adjust the current and charge densities to $j_{\rm m}$ and $\rho_{\rm GJ}$ simultaneously \citep[e.g., ][]{Me85}. In such cases, a significant accelerating electric field develops and causes the copious pair creation. The newly created pairs would screen the accelerating electric field for a temporary period of time \citep[e.g., ][]{S71, LMJL05, TA13}. The back-flowing particles from the outer acceleration region (OAR) modify the above description of the dynamics near the NSS. As a result of discharge at the OAR in the magnetosphere (e.g., outer gap or current sheet), some fraction of charged particles would come back to the NSS. Such back-flowing particles are actually seen in numerical studies \citep[e.g., ][]{H06, WS07, CB14, CPPS15}. The existence of the back-flow is favorable to explain the observed pulse profiles in the non-thermal soft $\gamma$-ray, X-ray and optical wavelengths, whose peaks are not aligned with the GeV $\gamma$-ray one \citep{TCS08, KK11, WTC13}. The back-flowing particles have been also considered to heat the NSS around the magnetic pole, whose signature is observed as the thermal pulsed emission in soft X-ray band \citep[e.g., ][]{ZP04}. The threshold of the occurrence of pair cascade near the NSS depends on the contribution of the back-flowing particles to the current and the charge densities \citep{B08}. The outflow from the NSS would also affect the dynamics in the OAR. The outgoing particles from the NSS contribute to the particle injection rate from the inner boundary of the OAR \citep[e.g., ][]{THSC06}. If an outflow from the NSS affects the accelerating electric field in the OAR, the current and number densities of the back-flowing particles change, and the resulting particle outflow from the NSS may be also modified. Through such non-linear interplay between the NSS and OAR, the global magnetosphere is expected to reach the steady or quasi-steady state (e.g., periodic behavior). \citet{Le+14} and \citet{TNC15} suggest that a non-stationary outer gap model is favored to reproduce sub-exponential cut-off feature in the GeV $\gamma$-ray spectrum observed with {\it Fermi}. In order to understand the global behavior of the magnetosphere, we need to link the dynamics of the region above the NSS and the OAR. In the first step, we focus on the dynamics of only the restricted region just above the NSS for given back-flowing particles. If a steady electric field just above the NSS exists, copious electron--positron pairs are produced via electromagnetic cascade. Such pairs flow into the OAR, and may screen out the electric field in the OAR. Therefore, the electric field just above the NSS should be almost screened out to activate the OAR. The local simulations have been performed to investigate the particle acceleration and the pair creation processes near the NSS \citep{BT07, L09, T10, B11, CB13, TA13, TH15}. Since the present global simulations are difficult to include the realistic pair-creation process with the actual mass-to-charge ratio \citep{SA02, WS07, WS11, YS12, PS14, CB14, CPPS15, B15, PSC15, PCTS15}, the local simulations are complementary. In order to link the local region above the NSS to the global magnetospheric structures, it is useful to model the properties of the outflowing particles from the NSS for arbitrary ratio $j_{\rm m}/j_{\rm GJ}$ and the back-flow from the OAR. \citet{T10}, \citet{TA13} and \citet{TH15} performed the local particle simulations to investigate the pair cascade near the NSS. In their regimes, a large number density of pairs are supplied to the OAR, so that the electric field in the OAR is screened by the copious pairs. Then, the back-flow from the OAR would be suppressed. In this context, the effect of the back-flowing particle has not been investigated in the local particle simulations so far. However, the OAR as a source of the back-flowing particles should exist if the pair cascade near the NSS fails to supply enough pairs. In this paper, we study the screening of the accelerating electric field above the NSS taking into account the effect of the back-flowing particles from the OAR. As we have mentioned above, we consider that the screening of the electric field near the NSS is a necessary condition to activate the OAR, because too much pair supply from the inner magnetosphere via a strong electric field would choke the OAR. The local condition of electric field screening near the NSS in the absence of the pair cascade is investigated for a given ratio, $j_{\rm m}/j_{\rm GJ}$, which is imposed in the global magnetospheric structure. In Section \ref{model}, we introduce our model with a particle outflow from the NSS, where the number density and momentum distribution of the back-flowing particles are given as model parameters. In Section \ref{analytic}, we analytically show the screening condition for the velocity of the plasma flow from the NSS in the case where the back-flowing particles are ultra-relativistic ($\gamma\gg1$). We see that the development of the accelerating electric field cannot be avoided for some combinations of the total current density and the contribution of the back-flowing particles. In Section \ref{numerical}, we introduce additional components of back-flowing particles, electron--positron pairs with a mildly relativistic temperature. Particle-in-Cell simulations are performed to investigate the screening conditions near the NSS. Implications of our results for the pulsar radio emission is discussed in Section \ref{radio}. We summarize our work in Section \ref{summary}.
% \label{summary} In order to activate the OAR, the electric field along the magnetic field just above the NSS should be screened out. In this paper, we investigate the condition on the electric field screening just above the NSS with back-flowing particles. We have focused on the case so-called anti-GJ condition, which would be established along the magnetic fields that connect to the OAR. Without the back-flows, the electric field cannot be screened out by the particles extracted from the NSS alone. First, we consider the case that particles accelerated in the OAR with Lorentz factor $\gamma\gg1$ (the beam component) are flowing back to the NSS [Figure \ref{image} (a)]. The current and charge densities of the beam component and particles from the NSS contribute to the total ones. Then, even in the anti-GJ case, there is some screened solutions in a certain region for the combination of the current densities $\alpha_{\rm m}$ and $\alpha_{\rm bk, bm}$ (Figure \ref{diagram}). We analytically introduce a parameter $\beta_{\rm ns}^{\rm req}$ to characterize the required flow speed of the particles extracted from the NSS. The electric field is not screened in the conditions so-called R-super-GJ and R-anti-GJ, which are defined by $\beta_{\rm ns}^{\rm req}$. However, we can expect that a quasi-thermal plasma with a mildly relativistic temperature screens out the electric field between the OAR and the outer boundary of the polar cap region [Figure \ref{image} (b)]. Some fraction of such thermal components can also flow back to the NSS. Using numerical simulations, we show that the thermal component can adjust both the current and charge densities to the required values. We obtain the minimum number density of the thermal component to screen the electric field in equations (\ref{screen1}) and (\ref{screen2}). We also investigate the ion-extracted case from the NSS. The difference of masses of electrons and ions causes bunches of particles, which are formed quasi-periodically. The period is linearly proportional to the length of the calculation domain. This may be important process for the coherent radio emission.
16
7
1607.02202
Recent γ-ray observations suggest that particle acceleration occurs at the outer region of the pulsar magnetosphere. The magnetic field lines in the outer acceleration region (OAR) are connected to the neutron star surface (NSS). If copious electron-positron pairs are produced near the NSS, such pairs flow into the OAR and screen the electric field there. To activate the OAR, the electromagnetic cascade due to the electric field near the NSS should be suppressed. However, since a return current is expected along the field lines through the OAR, the outflow extracted from the NSS alone cannot screen the electric field just above the NSS. In this paper, we analytically and numerically study the electric field screening at the NSS, taking into account the effects of the backflowing particles from the OAR. In certain limited cases, the electric field is screened without significant pair cascade if only ultra-relativistic particles (γ \gg 1) flow back to the NSS. On the other hand, if electron-positron pairs with a significant number density and mildly relativistic temperature, expected to distribute in a wide region of the magnetosphere, flow back to the NSS, these particles adjust the current and charge densities so that the electric field can be screened without pair cascade. We obtain the condition needed for the number density of particles to screen the electric field at the NSS. We also find that in the ion-extracted case from the NSS, bunches of particles are ejected to the outer region quasi-periodically, which is a possible mechanism of observed radio emission.
false
[ "NSS", "particle acceleration", "particles", "significant pair cascade", "pair cascade", "OAR", "observed radio emission", "such pairs", "the electric field", "only ultra-relativistic particles", "the NSS", "the outer acceleration region", "The magnetic field lines", "the backflowing particles", "the field lines", "the outer region", "account", "a significant number density", "these particles", "copious electron-positron pairs" ]
5.729302
3.596404
69
12408741
[ "Lingam, Manasvi", "Bhattacharjee, Amitava" ]
2016ApJ...829...51L
[ "Hall Current Effects in Mean-Field Dynamo Theory" ]
13
[ "Department of Astrophysical Sciences and Princeton Plasma Physics Laboratory, Princeton University, Princeton, NJ 08544, USA", "Department of Astrophysical Sciences and Princeton Plasma Physics Laboratory, Princeton University, Princeton, NJ 08544, USA; Max Planck-Princeton Center for Plasma Physics, Princeton University, Princeton, NJ 08544, USA" ]
[ "2016ApJ...829...87A", "2016PhPl...23k2104M", "2017NJPh...19a5007M", "2017PhPl...24b2107A", "2017PhPl...24f2307B", "2017PhRvE..96a3207K", "2018MNRAS.473.2822F", "2020JPlPh..86e8301L", "2020MNRAS.495.2771M", "2020PhLA..38426698B", "2021PhPl...28h2303M", "2021arXiv210407759M", "2023PhRvF...8e3701H" ]
[ "astronomy", "physics" ]
4
[ "dynamo", "galaxies: magnetic fields", "ISM: magnetic fields", "magnetohydrodynamics: MHD", "turbulence", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - Solar and Stellar Astrophysics", "Physics - Plasma Physics" ]
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[ "10.3847/0004-637X/829/1/51", "10.48550/arXiv.1607.05087" ]
1607
1607.05087_arXiv.txt
\label{SecIntro} The importance of large scale magnetic fields in the Universe, ranging from the cosmic to the planetary, cannot be overstated \citep{KZ08,Sub16}. They play a chief role in processes such as star formation, galaxy evolution, etc. and as a means of probing and interpreting astrophysical environments. The principal process responsible for generating these large scale fields is the dynamo mechanism, which has been extensively studied for nearly a century \citep{KR80,YIIY04,BS05}, but there are several key issues that are yet unresolved \citep{BSS12}. To date, most analyses in plasma astrophysics, and, by extension, in dynamos, rely on ideal or resistive magnetohydrodynamics (MHD) as the underlying physical model on account of its simplicity. However, the MHD approximation is \emph{not} universally valid, and there are many domains where other plasma effects become dominant. One of the simplest of these `extended' models is Hall MHD, in which the ions and electrons generally move with different flow velocities. This effect is manifested in the Ohm's law via the Hall term, and the relative magnitude of such a term is typically encoded in the Hall parameter - the ratio of the ion skin depth to the characteristic background scale length. In recent times, the Hall parameter has been shown to be `large' in diverse environments such as protoplanetary discs \citep{Ward07}, the ISM, the Earth's magnetotail \citep{Bhat04}, the solar corona and the solar wind \citep{Bhat04,GB07} to name a few. Furthermore, it is also likely to be highly relevant in space and laboratory settings \citep{Huba95,BMW01}, especially in experiments involving the Reversed Field Pinch \citep{Ji94,Ji99,Det04}. In each of these instances, understanding the dynamics of the magnetic fields would arguably necessitate the use of Hall MHD, as opposed to MHD, as the base physical model. Despite the ubiquity and importance of the Hall term, dynamo models that incorporate this term have been few and far between. The notable exceptions in large scale dynamo theory are the series of works undertaken by Mininni and collaborators \citep{MGM02,MGM03,MGM05,MAP07} as well as a few other publications \citep{Ji99,LM15}. There have also been some concomitant studies of small-scale Hall MHD dynamos, such as \citet{KR94,MHP03,GMD10}. The Hall MRI dynamo has been subjected to detailed investigations only in recent times \citep{SS02,KL13,LB16}. In this paper, we shall use incompressible Hall MHD and carry out an investigation along the lines of \citet{GD94} and \citet{GD95}, which will henceforth be referred to as GD94 and GD95 respectively. At this stage, a crucial comment pertaining to our approach is necessary. Our model does not include stochasticity and chaos \citep{VB97,BT07}, the shear-current effect \citep{RK03,RK04,SB16}, helicity flux contributions \citep{VC01,SB04,SSS07,EB14}, dynamical equations for the large- and small-scale helicities \citep{BF02,BSS12,Bla15}, and hyperresistivity \citep{BY95,BS05} to name a few. Although each and every one of these effects is undoubtedly important, the central goal of the paper is to understand how the Hall term acts in near-isolation. We shall show that the Hall term introduces several non-trivial effects, such as the existence of a non-positive definite diffusion coefficient, the drive towards a non-Alfv\'enic final state, and the possibility for dynamo action despite $\alpha$-quenching \citep{VC92,GD94,CH96} amongst others. Collectively, we argue that these non-trivial effects make a compelling case for carrying out in-depth Hall MHD dynamo analyses in the future.
\label{SecConc} In this paper, we have investigated the role of the Hall effect on conventional large scale dynamo theory. For the purposes of physical simplicity, we have used a very simple ansatz, along the lines of \citet{GD94,GD95}, as the basis of our investigations. As discussed in Sec. \ref{SecIntro}, the calculations presented here are only the first step in a more complete treatment of the Hall dynamo, which will be extended in future work(s) to include several effects that have been omitted for simplicity. Foremost among them is the effect of large-scale velocity shear, which has led to the prediction of the novel ``magnetic shear current effect'' \citep{SB16}. We have chosen to work with such a model as it helps clarify the effects engendered by the Hall term. In `real' systems, however, one would expect it to act in concert with all of the other phenomena alluded to in Sec. \ref{SecIntro}. Our final results are encapsulated by (\ref{betakin}), (\ref{alphaNkin}) and (\ref{alphaquench}). Thus, we find that the $\alpha$ and $\beta$ effects are modified when compared to their MHD counterparts, and that $\alpha$-quenching operates analogous to GD94 and GD95 as well as \citet{VC92,CH96}. Each of these results has important physical interpretations and consequences. We have also carried out a simple dimensional analysis in Sec. \ref{SSecHImp} and concluded that the dynamo coefficients are significantly altered when the Hall parameter is sufficiently large, and the latter does occur in several space, astrophysical and laboratory plasmas, as pointed out in Sec. \ref{SecIntro}. We observe that the null value of $\alpha$ is achieved via the process of `Beltramization' in Hall MHD, i.e. when the system attains a particular double Beltrami state, as opposed to `Alfv\'enization' \citep{DMV80,GFPL82} in MHD. We find that (i) wave turbulence (mediated via nonlinear Alfv\'en waves), (ii) relaxed states (the double Beltrami states can be thus interpreted), and (iii) large scale dynamo theory are intimately connected in Hall MHD. We believe that the connections between these three (highly important) areas of plasma physics certainly merit further investigation - a similar line of thought has also been advanced in a recent review by \citet{Moff16}, albeit for MHD. Moreover, there has been a great deal of interest in the diffusivity ($\beta$) becoming negative, and driving the large scale dynamo. Unlike many standard treatments that rely on large scale velocity shear or stochasticity \citep{VB97,RK03,SB16}, we have shown that the isotropic diffusion tensor can become negative under the influence of the Hall term, in concordance with earlier results \citep{MAP07}. We reiterate that the non-positive definiteness of $\beta$ induced by the Hall term is likely to complement (and supplement) the shear-current effect, and not necessarily supplant it. At this juncture, we proffer a potential explanation as to why $\beta < 0$ is possible. The backreactions of the Lorentz force $\left({\bf J} \times {\bf B}\right)$ are responsible for introducing the current helicity contribution to the $\alpha$ effect that is \emph{opposite} in sign to the kinematic fluid helicity \citep{PFL76,BS05}. However, the same term is also present in the Ohm's law of Hall MHD, and is responsible for new contributions to the dynamo coefficients - this is seen by inspecting the last two terms on the RHS of (\ref{EMFExp}). It is precisely these new contributions that destroy the positive definite nature of $\beta$. Hence, in a manner of speaking, the ${\bf J} \times {\bf B}$ term in the Ohm's law is solely responsible for generating new contributions that `oppose' the tendency for $\beta > 0$. Thus, it is quite plausible that the ${\bf J} \times {\bf B}$ term modifies \emph{both} the $\alpha$ and $\beta$ coefficients and enforces this `opposite' behaviour. As the Hall term is important in laboratory and astrophysical environments, gives rise to important and subtle physical effects in large scale dynamos, and connects the latter field with other fundamental areas of plasma physics, we suggest that further studies of the kind carried out in the present paper are timely.
16
7
1607.05087
The role of the Hall term on large-scale dynamo action is investigated by means of the first-order smoothing approximation. It is shown that the standard α coefficient is altered, and is zero when a specific double Beltrami state is attained, in contrast to the Alfvénic state for magnetohydrodynamical dynamos. The β coefficient is no longer positive definite, and thereby enables dynamo action even if α-quenching were to operate. The similarities and differences with the (magnetic) shear-current effect are pointed out, and a mechanism that may be potentially responsible for β \lt 0 is advanced. The results are compared against previous studies, and their astrophysical relevance is also highlighted.
false
[ "magnetohydrodynamical dynamos", "dynamo action", "α-quenching", "Alfvénic", "contrast", "a specific double Beltrami state", "large-scale dynamo action", "first", "previous studies", "means", "Beltrami", "the first-order smoothing approximation", "the Alfvénic state", "their astrophysical relevance", "the standard α coefficient", "Hall", "differences", "β \\lt", "the Hall term", "a mechanism" ]
11.615089
14.296091
2
12409227
[ "Pan, Margaret", "Nesvold, Erika R.", "Kuchner, Marc J." ]
2016ApJ...832...81P
[ "Apocenter Glow in Eccentric Debris Disks: Implications for Fomalhaut and ɛ Eridani" ]
50
[ "MIT Department of Earth, Atmospheric, and Planetary Sciences, Cambridge, MA 02139, USA ; Department of Astronomy and Astrophysics, University of Toronto, Toronto, Ontario M5S 3H4, Canada ; NASA Goddard Space Flight Center, Exoplanets and Stellar Astrophysics Laboratory, Greenbelt, MD 20771, USA", "Department of Applied Physics, University of Maryland Baltimore County, Baltimore, MD 21250, USA ; Department of Terrestrial Magnetism, Carnegie Institution of Washington, Washington, DC 20015, USA", "NASA Goddard Space Flight Center, Exoplanets and Stellar Astrophysics Laboratory, Greenbelt, MD 20771, USA" ]
[ "2017A&A...605A...7L", "2017AJ....154..225S", "2017ApJ...837L...6N", "2017ApJ...842....8M", "2017ApJ...842....9M", "2017ApJ...847...24T", "2017MNRAS.465.2595M", "2017MNRAS.469.3200B", "2017MNRAS.469.3518M", "2017MNRAS.470.3606H", "2018A&A...618A..38K", "2018ARA&A..56..541H", "2018ApJ...858..112L", "2018ApJ...861...72C", "2018ApJ...869...75M", "2018MNRAS.475.4924K", "2018MNRAS.479.5423M", "2018MNRAS.481...44F", "2018exha.book.....P", "2019A&A...626A..54M", "2019A&A...630A.142O", "2019A&A...631A.141S", "2019AJ....158..162F", "2019ApJ...872..184L", "2019ApJ...877L..32M", "2019MNRAS.484.1257M", "2020ApJ...898...55C", "2020MNRAS.497.1098H", "2020MNRAS.498.1319M", "2020RSOS....700063K", "2021ApJ...910...13S", "2021MNRAS.500.1604B", "2021MNRAS.503.4767P", "2022ApJ...925L...1M", "2022ApJ...932...23C", "2022ApJ...933L...1M", "2022ApJ...939...56F", "2022MNRAS.510.2538L", "2023A&A...669A...3S", "2023ASPC..534..539M", "2023ApJ...945..131S", "2023ApJ...954...14B", "2023MNRAS.525L..36L", "2023PSJ.....4...33P", "2023arXiv230110219P", "2023arXiv230317980F", "2024A&A...686A.206Y", "2024AJ....167..236G", "2024ApJ...961..245C", "2024MNRAS.529L.147L" ]
[ "astronomy" ]
6
[ "planetary systems", "planets and satellites: dynamical evolution and stability", "protoplanetary disks", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1994Icar..108...37R", "1999ApJ...527..918W", "2000ApJ...530..329T", "2000MNRAS.314..702D", "2001ApJ...554..778L", "2003A&A...411..559K", "2003ApJ...582.1141H", "2003ApJ...588.1110K", "2004ApJS..154..458S", "2005ApJ...620L..47M", "2005Natur.435.1067K", "2006ApJ...650..414S", "2006MNRAS.372L..14Q", "2009AJ....137...53S", "2009ApJ...693..734C", "2010A&A...518L.131E", "2011A&A...527A..57R", "2011ApJ...743L...6T", "2012A&A...539L...6R", "2012A&A...540A.125A", "2012A&A...546A..38L", "2012AJ....144...45K", "2012ApJ...754...57B", "2012ApJ...760L..17P", "2013ApJ...777..144N", "2014ApJ...780...65R", "2014ApJ...791L..11G", "2014MNRAS.443.2541P", "2015A&A...576A..72L", "2015ApJ...807L...7C", "2015ApJ...815...61N" ]
[ "10.3847/0004-637X/832/1/81", "10.48550/arXiv.1607.06798" ]
1607
1607.06798_arXiv.txt
\label{sec:introduction} More and more high-resolution images show that debris disks often take the form of rings, sometimes narrow, sometimes eccentric. Well-resolved examples include HR4796 \citep{schneider09, thalmann11, lagrange12}, Fomalhaut \citep{stapelfeldt04, kalas05}, HD 181327 \citep{schneider06}, $\zeta^2$ Reticuli \citep{eiroa10}, HD 202628 \citep{krist12}, and HD 115600 \citep{currie15}. These rings may indicate the presence of hidden planets, which can clear the central cavities in the rings and also excite the ring eccentricities via secular perturbations \citep{roques94, wyatt99, kuchner03, quillen06, chiang09, rodigas14, nesvold15}. At shorter wavelengths, regions of an eccentric disk near pericenter appear brighter because they receive more flux from the host star. \citet{wyatt99} named this phenomenon ``pericenter glow'' and developed a model for an eccentric debris ring interacting with a single planetary perturber. Their model disk consists of massless particles whose eccentricities differ only in the direction of the free eccentricity. The resulting ring suffices to explain the offset in the solar zodiacal cloud from the sun and to fit observations of several debris disks, yielding constraints on the disks' forced eccentricity and the masses of the hidden planetary perturbers. However, variations in the disk surface density also affect its apparent brightness, and in a steady-state disk the density should peak at apocenter simply because typical orbit velocities are slowest there. Though modeling done by \citet{wyatt99} predicted a brightness enhancement at pericenter for HR 4796 at 18.2~$\mu$m, their disk model showed a 2\% density enhancement at apocenter. Analogous apocentric density enhancements occur in more recent dynamical models of eccentric planets interacting with disks \citep[see, for example,][]{nesvold13, pearce14}. Indeed, submillimeter observations of the very well-observed eccentric Fomalhaut disk consistently suggest apocentric brightness enhancements. JCMT images of Fomalhaut by \citet{holland03} show slight enhancements of the flux near apocenter; these enhancements are less than the quoted uncertainty in the photometry, but they appear in both 450~$\mu$m and 850~$\mu$m bands. When \citet{marsh05} imaged the Fomalhaut disk at 350~$\mu$m using the SHARC II (Submillimeter High Angular Resolution Camera II) at the Caltech Submillimeter Observatory, they found that the ring has an apocentric enhancement of approximately 14\% in integrated column density. More recently, \citet{ricci12} imaged Fomalhaut's disk at 7~mm with the ATCA and noted that the lobe of the disk near apocenter ``...appears to be more extended, showing two possible asymmetric structures toward east and south.'' The highest resolution ALMA images of Fomalhaut by \citet{boley12} at 350~GHz (1~mm) also show enhanced flux near apocenter in the maps corrected for the single-dish beam. Here we describe a new model for debris rings that illustrates the wavelength dependence in the apocenter/pericenter flux ratio due to the competing effects of azimuthal asymmetries in dust density and stellar illumination. Our primary interest here is mid-infrared and longer wavelengths, so we focus on dust particles and planetesimals large enough to avoid radiation pressure effects and consider only absorbed and re-radiated, rather than scattered, emission. In Section \ref{sec:semianalytic}, we begin by describing a semi-analytic model for estimating the surface density of a steady-state distribution of collisionless planetesimals and show that the density of a dust ring varies with longitude and peaks at apocenter. In Section \ref{sec:smack}, we verify this result for a collisional ring using SMACK \citep{nesvold13}, a numerical model of debris disk evolution that incorporates both collisions and dynamics in 3D. Finally in Section \ref{sec:rerad}, we combine a simple dust reradiation model with our models of surface density to simulate the brightness of the Fomalhaut and $\epsilon$ Eridani rings, and demonstrate that the ratio between pericenter and apocenter flux varies with wavelength. We summarize our results and discuss the implications for future observations of eccentric rings in Section \ref{sec:conclusions}.
\label{sec:conclusions} Using both semi-analytic and numerical modeling of the azimuthal dust distribution in an eccentric ring of colliding planetesimals, we have studied the wavelength dependence of surface brightness variations using simple assumptions about dust radiative properties and size distributions. We argued that several far-infrared and (sub)millimeter images of Fomalhaut and $\epsilon$ Eridani obtained with Herschel, JCMT, SHARC II, ALMA, and ATCA \citep{holland03,marsh05,ricci12,boley12,greaves14} should be reinterpreted as suggestions or examples of apocenter glow. This reinterpretation also yields new constraints on the grain properties and size distributions from the existing data. Our modeling work also has implications for future observations of debris disks. The James Webb Space Telescope will be a powerful new source of debris disk images, observing at 5-28 microns with the MIRI instrument. Figure~\ref{fig:fluxampsmalldisk} illustrates that this wavelength range is particularly sensitive to pericenter glow. We predict apocenter/pericenter flux ratios as small as 0.1 in this range for dust emitting via the highly temperature-sensitive Wien law. {\em ALMA}, on the other hand, will operate at wavelengths from 3mm to 400$\mu$m, primarily continuing to detect apocenter glow. {\em ALMA} images will be especially useful because, as Figures~\ref{fig:fluxampfomalhaut}, \ref{fig:fluxampepseri}, \ref{fig:fluxampsmalldisk} show, the apocenter/pericenter flux ratio becomes insensitive to dust properties at millimeter wavelengths. For a fixed disk mass, changes in $q$ and $\beta$ mostly affect the flux contributed by the smallest particles: increasing $q$ increases the number of very small particles, while increasing $\beta$ decreases the flux reradiated at wavelengths larger than the particle size. However, for the longest wavelengths, $\lambda\gg s$, the flux from small particles becomes negligible. To lowest order the millimeter apocenter/pericenter flux ratio therefore depends only on the apocenter/pericenter temperature and density contrasts. The largest bodies radiate efficiently and have effective temperature \begin{equation} T\sim T_* (R_*/r)^{1/2}\propto 1\pm e/2 \end{equation} where we assume the star is a blackbody with effective temperature $T_*$ and radius $R_*$ and, in the last step, apply Eq.~\ref{eqn:r} with $f=0$ (top sign) for pericenter and $f=\pi$ for apocenter (bottom sign). Together with Eq.~\ref{eqn:ell} evaluated at $f=0$ and $\pi$, this gives \begin{equation} \mathrm{apocenter/pericenter\;flux\;ratio} \simeq \frac{1-e/2}{1+e/2}\frac{1+e}{1-e} \simeq 1+e \end{equation} where, as before, we take only lowest order terms in the eccentricity. The millimeter apocenter/pericenter flux ratio thus provides a direct estimate of the disk eccentricity. Some systems, like $\epsilon$ Eridani, no doubt contain additional structures that will complicate interpretation of their images. However, with this new understanding of apocenter glow and its wavelength dependence, we can begin future studies of debris disk images pointed in the right direction.
16
7
1607.06798
Debris disks often take the form of eccentric rings with azimuthal asymmetries in surface brightness. Such disks are often described as showing “pericenter glow,” an enhancement of the disk brightness in regions nearest the central star. At long wavelengths, however, the disk apocenters should appear brighter than their pericenters: in the long-wavelength limit, we find that the apocenter/pericenter flux ratio scales as 1+e for disk eccentricity e. We produce new models of this “apocenter glow” to explore its causes and wavelength dependence and study its potential as a probe of dust grain properties. Based on our models, we argue that several far-infrared and (sub)millimeter images of the Fomalhaut and ɛ Eridani debris rings obtained with Herschel, JCMT, SHARC II, ALMA, and ATCA should be reinterpreted as suggestions or examples of apocenter glow. This reinterpretation yields new constraints on the disks’ dust grain properties and size distributions.
false
[ "size distributions", "dust grain properties", "disk eccentricity", "Debris disks", "surface brightness", "ɛ Eridani debris rings", "pericenter glow", "apocenter", "Such disks", "long wavelengths", "eccentric rings", "SHARC II", "azimuthal asymmetries", "the disk apocenters", "examples", "regions", "suggestions", "the disk brightness", "ATCA", "ALMA" ]
9.202697
13.342753
-1
12473231
[ "Bulla, M.", "Sim, S. A.", "Kromer, M.", "Seitenzahl, I. R.", "Fink, M.", "Ciaraldi-Schoolmann, F.", "Röpke, F. K.", "Hillebrandt, W.", "Pakmor, R.", "Ruiter, A. J.", "Taubenberger, S." ]
2016MNRAS.462.1039B
[ "Predicting polarization signatures for double-detonation and delayed-detonation models of Type Ia supernovae" ]
36
[ "Astrophysics Research Centre, School of Mathematics and Physics, Queen's University Belfast, Belfast BT7 1NN, UK", "Astrophysics Research Centre, School of Mathematics and Physics, Queen's University Belfast, Belfast BT7 1NN, UK; ARC Centre of Excellence for All-sky Astrophysics (CAASTRO), Australian National University, Canberra, ACT 2611, Australia", "Department of Astronomy, The Oskar Klein Centre, Stockholm University, AlbaNova, SE-106 91 Stockholm, Sweden", "ARC Centre of Excellence for All-sky Astrophysics (CAASTRO), Australian National University, Canberra, ACT 2611, Australia; Research School of Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611, Australia", "Institut für Theoretische Physik und Astrophysik, Universität Würzburg, Emil-Fischer-Straße 31, D-97074 Würzburg, Germany", "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, D-85748 Garching bei München, Germany", "Heidelberger Institut für Theoretische Studien, Schloss-Wolfsbrunnenweg 35, D-69118 Heidelberg, Germany; Zentrum für Astronomie der Universität Heidelberg, Institut für Theoretische Astrophysik, Philosophenweg 12, D-69120 Heidelberg, Germany", "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, D-85748 Garching bei München, Germany", "Heidelberger Institut für Theoretische Studien, Schloss-Wolfsbrunnenweg 35, D-69118 Heidelberg, Germany", "ARC Centre of Excellence for All-sky Astrophysics (CAASTRO), Australian National University, Canberra, ACT 2611, Australia; Research School of Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611, Australia", "Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, D-85748 Garching bei München, Germany; European Southern Observatory, Karl-Schwarzschild-Str. 2, D-85748 Garching, Germany" ]
[ "2017A&A...603A.136P", "2017ApJ...841...62M", "2017PhDT.........1B", "2017hsn..book..769S", "2018ApJ...862...27G", "2018MNRAS.477.3567M", "2018MNRAS.479L..70D", "2018PhR...736....1L", "2018SSRv..214...72R", "2019ApJ...878...86A", "2019MNRAS.486.5528L", "2019MNRAS.489.5037B", "2019MNRAS.490..578C", "2019NatAs...3...99B", "2019NewAR..8701535S", "2019supe.book...69R", "2020A&A...644A.162S", "2020ApJ...897..159P", "2020ApJ...902...46Y", "2020MNRAS.494.5811L", "2021ApJ...906...93F", "2021ApJ...906...99L", "2021ApJ...909..152L", "2021ApJ...914...50T", "2021ApJ...919..126B", "2021ApJ...922..186H", "2021MNRAS.503.4667Z", "2022ApJ...939...18Y", "2022MNRAS.509.4058P", "2022MNRAS.511.2994L", "2022MNRAS.511.3682G", "2023ApJ...958..178L", "2023MNRAS.520..560H", "2023RAA....23h2001L", "2024ApJ...962..125C", "2024MNRAS.528.3875M" ]
[ "astronomy" ]
2
[ "hydrodynamics", "polarization", "radiative transfer", "methods: numerical", "supernovae: general", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1947ApJ...106..457H", "1980ApJ...237..142T", "1980SSRv...27..563N", "1989ApJS...71..951J", "1990ApJ...354L..53L", "1990ApJ...361..244L", "1991A&A...245..114K", "1991ApJ...375..264J", "1995ApJ...444..831H", "1996ApJ...457..500H", "1997ApJ...475..740N", "1997ApJ...476L..27W", "1998AJ....116.1009R", "1999ApJ...517..565P", "2001ApJ...553..861L", "2001ApJ...556..302H", "2003ApJ...591.1110W", "2003ApJ...593..788K", "2004ApJ...607..391G", "2004ApJ...610..876K", "2004ApJ...612L..37P", "2005ApJ...623L..37M", "2005ApJ...632..450L", "2005MNRAS.357..200M", "2006A&A...460..793S", "2006ApJ...645..470T", "2006ApJ...653..490W", "2006NewAR..50..470H", "2006PASP..118..146P", "2007A&A...464..683R", "2007A&A...476.1133F", "2007ApJ...660.1344R", "2007ApJ...662L..95B", "2007MNRAS.375..154S", "2007Sci...315..212W", "2008ARA&A..46..433W", "2008ApJ...683L.127H", "2009A&A...508..229P", "2009ApJ...699.1365S", "2009MNRAS.398.1809K", "2009Natur.460..869K", "2010A&A...514A..53F", "2010A&A...515A..89M", "2010ApJ...719.1067K", "2011MNRAS.417..408R", "2012A&A...545A...7P", "2012ApJ...747L..10P", "2012ApJ...750L..19R", "2012MNRAS.427..994G", "2013A&A...554A..67S", "2013A&A...559A.117C", "2013ApJ...774..137M", "2013FrPhy...8..116H", "2013MNRAS.429.1156S", "2013MNRAS.429.2127B", "2013MNRAS.433L..20M", "2013MNRAS.436..333S", "2014ATel.5830....1P", "2014MNRAS.437..338C", "2014MNRAS.440.1498S", "2014MNRAS.444.3258M", "2014MNRAS.445.2535S", "2015ApJ...814L...2F", "2015ApJS..220...20Z", "2015MNRAS.448.2766B", "2015MNRAS.449.1441K", "2015MNRAS.450..967B", "2015MNRAS.451.1973S", "2015MNRAS.454.3816C", "2016ApJ...817...53S", "2016ApJ...828...24P", "2016MNRAS.455.1060B" ]
[ "10.1093/mnras/stw1733", "10.48550/arXiv.1607.04081" ]
1607
1607.04081_arXiv.txt
\label{introduction} Despite their relevance to cosmology \citep{riess1998,perlmutter1999}, Type Ia supernovae (SNe~Ia) are still poorly understood. While believed to stem from thermonuclear explosions of carbon-oxygen white dwarfs (WDs), answers to the questions of when, why and how these events are triggered remain unclear \citep{hillebrandt2013,maoz2014}. Recent explosion simulations have led to a better understanding of the physics involved, and comparisons of synthetic light curves and spectra with observations have played a key role in identifying which explosions scenarios are most promising. However, unambiguous discrimination between models is still challenging, even for the best observed nearby supernovae \citep{roepke2012}. Polarization offers a unique opportunity to discriminate between the variety of possible explosion scenarios. The observational evidence that SNe~Ia are associated with rather low levels of polarization ($\lesssim$~1~per~cent) demands modest asphericities in the progenitor system and/or explosion mechanism, thus providing the means to effectively test different explosion models. Although predictions have been made using idealized geometries with simple departures from spherical symmetry -- such as ellipsoidal structures \citep{wang1997,howell2001,hoeflich2006,patat2012}, clumped and toroidal shells \citep{kasen2003} or large scale asymmetries associated with ejecta overrunning the companion star \citep{kasen2004} -- until recently, no spectropolarimetric studies had been made for full multi-dimensional hydrodynamic explosion simulations. We have therefore started our \revised{long-term} project aiming to predict polarization signatures for \revised{a series of} modern SN~Ia hydrodynamic models. \revised{In the first paper of this series} \citep{bulla2016}, \revised{we have demonstrated the power of this approach for} the violent merger of two 1.1 and 0.9~M$_{\odot}$ WDs as presented by \citet{pakmor2012}. Despite matching luminosities and spectra of normal SNe~Ia reasonably well, this model was found to be too asymmetric to reproduce the polarization levels seen for the majority of normal SNe~Ia. \revised{In this second paper}, we carry out polarization spectral synthesis for examples of two alternative scenarios: the ``delayed-detonation" of a near Chandrasekhar mass (\mbox{M$_\text{ch}$}) WD and the ``double-detonation" of a sub-\mbox{M$_\text{ch}$} WD. In the ``delayed-detonation" model \citep[e.g.][]{khokhlov1991,hoeflich1995,plewa2004,roepke2007b,kasen2009,blondin2013,seitenzahl2013}, a WD is thought to accrete material from a non-degenerate companion star and explode following carbon ignition near the WD centre which occurs when the mass has grown close to \mbox{M$_\text{ch}$}. Initially, a carbon deflagration is ignited but delayed-detonation models posit that during the evolution of the explosion a detonation occurs. The second scenario we consider, the ``double-detonation" model \citep[e.g.][]{nomoto1980,taam1980,livne1990a,shen2009,fink2010,moll2013}, describes the explosion of a sub-\mbox{M$_\text{ch}$} WD. In this model, the explosion is triggered by the detonation of a helium surface layer that has been accreted from a helium-rich companion star. The shock wave from this detonation then triggers a second detonation in the core. Both delayed-detonation and double-detonation scenarios have received considerable attention over the past decade as they reproduce SN~Ia light curves and spectra reasonably well \citep{hoeflich1995,hoeflich1996,kasen2009,kromer2010,roepke2012,sim2013,blondin2015} and could provide an important contribution to the SN~Ia population \citep[e.g.][]{hachisu2008,mennekens2010,ruiter2011}. In this paper we perform polarization calculations for one explosion model of \citet{fink2010} and one model of \citet{seitenzahl2013}. The former is chosen as it has been used in several studies (\citealt{scalzo2014a}a; \citealt{scalzo2014b}b; \citealt{childress2015}; \citealt{kosenko2015}) as a benchmark for the double-detonation scenario. The latter is selected since it has been widely used (\citealt{roepke2012}; \citealt{summa2013}; \citealt{scalzo2014a}a; \citealt{childress2015}; \citealt{fransson2015}; \citealt{kosenko2015}; \citealt{sneden2016}) as representative of a class of delayed-detonation models in which a deflagration-to-detonation transition (DDT) is assumed to occur spontaneously during the propagation of the deflagration. In Section~\ref{models} we summarize properties of the two specific explosion models, while in Section~\ref{simul} we discuss the details of our radiative transfer calculations. We present synthetic observables for both models in Section~\ref{synthobs} and comparisons with spectropolarimetric data of SNe~Ia in Section~\ref{compobs}. Finally, we discuss our results and draw conclusions in Section~\ref{conclusions}.
\label{conclusions} We have presented polarization calculations for one double-detonation model from \citet[][here referred to as D-DET]{fink2010} and one delayed-detonation model from \citet[][here N100-DDT]{seitenzahl2013}. Using the techniques of \citet{bulla2015}, we calculated polarization spectra around maximum light and focused on three orientations for the 2D \subch model and five for the 3D \ch model. In order to map out the range of polarization covered by the models, we also performed extra-calculations that provide lower signal-to-noise spectra for an additional 17 (\mbox{D-DET}) and 21 (\mbox{N100-DDT}) viewing angles. For both models, the overall polarization levels are low at all the epochs considered in this study (between 10 and 30~d after explosion): they are consistently below 1~per~cent, with the largest values generally observed for the \ch model. While the 2D \subch model is axi-symmetric by construction, polarization spectra for the 3D \ch model verify that polarimetry could be used to identify deviations from a single-axis geometry. A common behaviour between the two explosion models is the relative rise in polarization level from blue to red wavelengths. Our calculations confirm that this effect, which is regularly observed in both normal \citep[e.g.][]{wang1997} and sub-luminous \citep[e.g.][]{howell2001} SNe~Ia, can be ascribed to the decreasing line blanketing when moving to longer wavelengths, and is fairly well captured in the explosion models: in our simulations around maximum light, the fraction of escaping packets having a depolarizing line interaction as last event drops from about 0.8 at 4000~\AA{} to only 0.3 at 7000~\AA. Polarization spectra predicted by our models match particularly well those observed for the majority of normal SNe~Ia, as clearly highlighted by the comparison with the two well-studied supernovae SN~2001el and SN~2012fr. Polarization levels in the pseudo-continuum range (between 6500 and 7500~\AA) are found to decrease from $\sim$~0.1$-$0.3~per~cent at maximum light to lower values at later times, in strikingly good agreement with observations. Higher degrees of polarization are predicted at wavelengths corresponding to the troughs of absorption lines. In particular, the set of polarization signals extracted across the Si\,{\sc ii}~$\lambda6355$ feature ($\lesssim$~1~per~cent) provides a remarkable match with the distribution of values observed in normal SNe~Ia \citep{wang2007,patat2009,maund2013}, while the low degrees of polarization found in the O\,{\sc i}~$\lambda7774$ region are consistent with the non-detection of this feature in currently available data. Taken together, these findings lead us to conclude that the geometries of both the explosion models considered here are likely to be broadly consistent with those of normal SNe~Ia. However, none of our models is able to reproduce the high-velocity ($\sim$~15\,000$-$20\,000~km~s$^{-1}$) components of the Ca\,{\sc ii} IR triplet and Si\,{\sc ii}~$\lambda6355$ lines that are frequently observed in both intensity (e.g. \citealt{gerardy2004}; \citealt{mazzali2005a}a; \citealt{childress2014,maguire2014,silverman2015,zhao2015}) and polarization \citep[e.g.][]{wang2003,leonard2005,wang2006,maund2013} spectra of SNe~Ia. Although the \subch scenario is more promising in this sense -- as the calcium in the outer shell of this model leads to velocities in the right range for some orientations -- none of our models is able to simultaneously produce the spectroscopic and polarimetric signatures characteristic of this high-velocity material. The failure to predict silicon and calcium in the right location and with the right geometry clearly indicates that more work is needed to understand whether the two scenarios investigated in this study can account for the high-velocity features observed in the Si\,{\sc ii}~$\lambda6355$ and Ca\,{\sc ii} IR triplet lines. The two models presented in this paper are hard to distinguish based on their polarization signatures. In both, polarizing electron scattering contributions originate in rather spherical regions of the ejecta (between $\sim$~5000 and 15\,000~km~s$^{-1}$), thus leading to polarization levels that are typically low and similar for the two models at most wavelengths. It is important to note, however, that we have only calculated polarization signatures for two specific models and the results presented here are not necessarily representative of the whole \citet{fink2010} double-detonation and \citet{seitenzahl2013} DDT studies. While the different double-detonation models of \citet{fink2010} are all characterized by relatively similar geometries -- and are thus expected to produce comparable polarization signals -- the DDT models of \citet{seitenzahl2013} yield ejecta whose morphologies are strongly dependent on the ignition configuration (see also Section~\ref{chmodel}). In particular, compared to the N100 model used in this paper, models in which the explosion is initiated with a smaller number of kernels (1 to 40) are typically more asymmetric and thus likely to be more polarized. In the future, it will be crucial to map out the polarization range spanned by the full set of models of \citet{fink2010} and \citet{seitenzahl2013} to understand whether the two scenarios can be effectively distinguished via variations between the polarization signatures across the model sequences. This study is part of a long-term project aimed to test multi-dimensional hydrodynamic explosion models via their polarization signatures. In a previous paper \citep{bulla2016}, we have demonstrated that the violent merger of a 1.1 and a 0.9~M$_{\odot}$ WD \citep{pakmor2012} is unlikely to explain polarization features observed for the bulk of normal SNe~Ia. In contrast, this work shows that the realizations of the double-detonation and DDT scenarios studied here produce polarization levels in much better agreement with spectropolarimetric data of normal SNe~Ia. To investigate whether the polarization properties of the discussed model classes are representative for the observed SN~Ia population, in the future we aim to map out the parameter space of all the three explosion scenarios.
16
7
1607.04081
Calculations of synthetic spectropolarimetry are one means to test multidimensional explosion models for Type Ia supernovae. In a recent paper, we demonstrated that the violent merger of a 1.1 and 0.9 M<SUB>⊙</SUB> white dwarf binary system is too asymmetric to explain the low polarization levels commonly observed in normal Type Ia supernovae. Here, we present polarization simulations for two alternative scenarios: the sub-Chandrasekhar mass double-detonation and the Chandrasekhar mass delayed-detonation model. Specifically, we study a 2D double-detonation model and a 3D delayed-detonation model, and calculate polarization spectra for multiple observer orientations in both cases. We find modest polarization levels (&lt;1 per cent) for both explosion models. Polarization in the continuum peaks at ∼0.1-0.3 per cent and decreases after maximum light, in excellent agreement with spectropolarimetric data of normal Type Ia supernovae. Higher degrees of polarization are found across individual spectral lines. In particular, the synthetic Si II λ6355 profiles are polarized at levels that match remarkably well the values observed in normal Type Ia supernovae, while the low degrees of polarization predicted across the O I λ7774 region are consistent with the non-detection of this feature in current data. We conclude that our models can reproduce many of the characteristics of both flux and polarization spectra for well-studied Type Ia supernovae, such as SN 2001el and SN 2012fr. However, the two models considered here cannot account for the unusually high level of polarization observed in extreme cases such as SN 2004dt.
false
[ "normal Type Ia supernovae", "Type Ia supernovae", "SN 2012fr", "SN 2004dt", "SN 2001el", "polarization spectra", "modest polarization levels", "SN", "Polarization", "polarization", "polarization simulations", "multidimensional explosion models", "spectropolarimetric data", "current data", "extreme cases", "the low polarization levels", "levels", "individual spectral lines", "Higher degrees", "multiple observer orientations" ]
14.076636
-0.531136
0
12464905
[ "Babichev, Eugeny", "Marzola, Luca", "Raidal, Martti", "Schmidt-May, Angnis", "Urban, Federico", "Veermäe, Hardi", "von Strauss, Mikael" ]
2016JCAP...09..016B
[ "Heavy spin-2 Dark Matter" ]
111
[ "Laboratoire de Physique Théorique, CNRS, Univ. Paris-Sud, Université Paris-Saclay, 91405 Orsay, France ; UPMC-CNRS, UMR7095, Institut d'Astrophysique de Paris, GReCO, 98bis boulevard Arago, F-75014 Paris, France;", "National Institute of Chemical Physics and Biophysics, Rävala 10, 10143 Tallinn, Estonia ; Laboratory of Theoretical Physics, Institute of Physics, University of Tartu, Ravila 14c, 50411 Tartu, Estonia;", "National Institute of Chemical Physics and Biophysics, Rävala 10, 10143 Tallinn, Estonia ; Laboratory of Theoretical Physics, Institute of Physics, University of Tartu, Ravila 14c, 50411 Tartu, Estonia;", "Institut für Theoretische Physik, Eidgenössische Technische Hochschule Zürich, Wolfgang-Pauli-Strasse 27, 8093 Zürich, Switzerland", "National Institute of Chemical Physics and Biophysics, Rävala 10, 10143 Tallinn, Estonia", "National Institute of Chemical Physics and Biophysics, Rävala 10, 10143 Tallinn, Estonia", "UPMC-CNRS, UMR7095, Institut d'Astrophysique de Paris, GReCO, 98bis boulevard Arago, F-75014 Paris, France" ]
[ "2016PhRvD..94h4055B", "2016arXiv161001016K", "2017AN....338.1034U", "2017APh....93...56M", "2017IJMPA..3230023B", "2017JHEP...06..130G", "2017JHEP...09..104C", "2017PhLB..774..676T", "2017PhRvD..95d4040K", "2017PhRvD..95l4049H", "2017PhRvD..96f4003T", "2017PhRvD..96j3519C", "2017PhRvD..96j4039A", "2017PhRvD..96l4023M", "2017PhRvL.119k1101M", "2017PhRvL.119y1301B", "2018ForPh..6600031L", "2018IJMPA..3345010U", "2018JCAP...01..014G", "2018JCAP...02..027G", "2018JCAP...07..012M", "2018JCAP...08..016D", "2018JHEP...09..044G", "2018JHEP...09..135E", "2018PDU....22...67V", "2018PhLB..785..159F", "2018PhRvD..97b4010M", "2018PhRvD..97d4002A", "2018PhRvD..97d4005A", "2018PhRvD..98j4014N", "2018Univ....4...90U", "2018arXiv181208686L", "2019CQGra..36n3001B", "2019CQGra..36w5010T", "2019JCAP...01..021G", "2019JMP....60d2501S", "2019PhRvD..99l3011N", "2019arXiv190708010V", "2020CQGra..37b4001D", "2020CQGra..37b5013T", "2020CQGra..37t5020Y", "2020IJMPA..3550120Y", "2020JCAP...04..046H", "2020JCAP...09..024L", "2020JHEP...01..131A", "2020JHEP...07..121A", "2020JHEP...07..231L", "2020MPLA...3550330S", "2020PDU....2700452E", "2020PhR...842....1A", "2020PhRvD.101h4026D", "2020PhRvD.102l3529L", "2020PhRvD.102l4049A", "2020PhRvD.102l5031C", "2020PhRvL.124u1101B", "2021JCAP...04..053A", "2021JCAP...05..001H", "2021JCAP...05..007U", "2021JCAP...05..065K", "2021JCAP...09..035C", "2021JCAP...11..001H", "2021JHEP...01..089R", "2021JHEP...04..217W", "2021JHEP...05..254C", "2021JHEP...11..036D", "2021JHEP...12..220L", "2021NatSR..1122528M", "2021PhRvD.104b4025D", "2021PhRvD.104e5046C", "2021PhRvD.104j4049Y", "2021ScPP...10..101F", "2021arXiv210609030H", "2022AnP...53400224A", "2022MNRAS.515.5646D", "2022PhRvD.105a6016C", "2022PhRvD.105e6019J", "2022PhRvD.105i6037Z", "2022PhRvD.106d3021C", "2022PhRvL.128h1806C", "2022arXiv220313720W", "2023EPJC...83.1120B", "2023JCAP...02..022H", "2023JCAP...09..021W", "2023JCAP...11..081G", "2023JHEP...05..181K", "2023JHEP...08..193L", "2023JHEP...09..101D", "2023NJPh...25c3019S", "2023PhLB..84738280B", "2023PhRvD.107j4007M", "2023PhRvD.107j4012G", "2023PhRvD.108d1502D", "2023PhRvD.108j4006O", "2023PhRvD.108j4023G", "2023PhRvD.108l4048E", "2023PhRvE.108e5305J", "2023arXiv230108267U", "2023arXiv230907904C", "2023arXiv231016007A", "2023arXiv231209042K", "2024GReGr..56...46P", "2024GReGr..56...64N", "2024JCAP...03..016G", "2024JCAP...04..053G", "2024JHEP...05..315D", "2024PhRvD.109i5012M", "2024PhRvD.109i5035C", "2024PhRvD.109i6001C", "2024PhRvD.109l4006W", "2024arXiv240100043Z", "2024arXiv240203984C" ]
[ "astronomy", "physics" ]
12
[ "High Energy Physics - Theory", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology" ]
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[ "10.1088/1475-7516/2016/09/016", "10.48550/arXiv.1607.03497" ]
1607
1607.03497_arXiv.txt
\label{sec:intro} Numerous cosmological and astrophysical observations have confirmed the presence of a Dark Matter (DM) component in our Universe. Until now this unknown type of matter has been seen only through its gravitational interactions, which resemble those of ordinary matter. Its effects are visible in the rotation curves and velocity dispersions of galaxies, gravitational lensing, matter distribution power spectra, structure formation, Baryon Acoustic Oscillations and angular power spectrum of the Cosmic Microwave Background~\cite{Agashe:2014kda}. The standard paradigm treats the unknown DM particle as a cold relic density which has been created through a model-dependent production mechanism in the early Universe. General Relativity (GR) as the theory for gravity (including a cosmological constant $\Lambda$ which accounts for the observed amount of Dark Energy) together with a particle physics model for cold Dark Matter (CDM) yield the concordance description of cosmology, the $\Lambda$CDM model. See \cite{Bull:2015stt} for a recent review of its status quo. The most popular DM models moreover assume that the DM particle is weakly coupled to baryonic matter and hence might be produced in colliders, directly detected in dedicated experiments or indirectly observed through astro-particle signatures. From a theoretical perspective, many of these models lose some of their attractiveness because, typically, they are either not very well motivated from fundamental principles or they introduce a large number of unobserved additional fields (such as Supersymmetry). Unfortunately, on the experimental side, all attempts to produce or detect the DM particle have remained unsuccessful so far~\cite{Agashe:2014kda, Ackermann:2015zua, Ahnen:2016qkx, Ackermann:2015lka, Khachatryan:2014rra, Gaskins:2016cha}. The absence of any signatures for DM apart from its gravitational effects motivates a shift of paradigm in the way to think about the nature of DM. Instead of augmenting the Standard Model (SM) by an additional field, we suggest that the DM particle may instead arise in a minimal extension of the gravitational sector, namely in the form of an additional massive spin-2 field. To us this seems to be a natural and well-motivated proposal, since there is no evidence supporting the fact that DM shares the quantum numbers of one of the SM particles and we only observe it through its gravitational interactions. General Relativity can be treated as the unique theory of a single massless spin-2 particle, the graviton. We will colloquially refer to this point of view as the standard description of gravity. Since massless spin-2 fields cannot interact with each other~\cite{Boulanger:2000rq}, the most natural and minimal addition to gravity is that of a massive spin-2 field. Studying the effects of a massive spin-2 field in addition to standard gravity amounts to answering fundamental questions of field theory. Remarkably, for several decades it was believed that no consistent theory for gravitating massive spin-2 fields can be formulated owing to the unavoidable presence of a fatal ghost instability~\cite{Boulware:1973my}. Only recently has the unique description which avoids the ghost been found~\cite{deRham:2010kj, Hassan:2011hr, Hassan:2011tf, Hassan:2011zd}. Since it involves an additional dynamical tensor field, which mixes with the gravitational metric, the corresponding theory has been titled ``bimetric theory of gravity". If fundamental massive spin-2 particles exist, they are described by this unique theory, which automatically leads to a modification of gravity. For the history and detailed reviews of theories for massive spin-2 fields we refer the reader to~\cite{Hinterbichler:2011tt, deRham:2014zqa, Schmidt-May:2015vnx}. Following ideas outlined in \cite{Schmidt-May:2015vnx, Schmidt-May:2016hsx}, it has been proposed that the existence of a massive spin-2 particle can explain all the effects related to DM~\cite{Babichev:2016hir} (see also \cite{Aoki:2016zgp}). The present work is dedicated to providing details of and more insights into this novel proposal. \paragraph{Summary of results.} Being a modification of gravity, bimetric theory must satisfy constraints coming from Solar System tests and cosmology. We confirm that a large value for the spin-2 mass, together with a small value for the ``second Planck mass" of the metric that does not couple directly to matter, imply that the static spherically symmetric and cosmological solutions of bimetric theory always resemble those of GR. In this parameter region, where bimetric theory passes all observational tests of GR, the additional massive spin-2 field continues to gravitate but decouples from matter, automatically providing an ideal DM candidate. We derive the conditions which ensure the validity of a perturbative treatment of bimetric theory for the interesting parameter regions and energy regimes. The structure of cubic and higher interactions for the spin-2 fields forbids a decay of the massive field into massless gravitons, resulting in a discriminating feature of the bimetric model. Requiring sufficient spin-2 DM to be produced in the early Universe and imposing constraints coming from its possible decay into SM fields, we obtain the allowed region in the bimetric parameter space: The spin-2 mass has to lie within the narrow region of $1~\TeV\lesssim\mfp\lesssim66~\TeV$ and the ratio of Planck masses must satisfy $10^{-11}\lesssim\alpha\lesssim10^{-15}$. This region overlaps with the one where classical solutions to the bimetric equations resemble GR to a very high precision and therefore the theory passes all observational tests. Moreover, our setup introduces no additional energy scale significantly higher than the weak scale and thus does not create any new hierarchy problems with respect to GR. Our novel DM proposal naturally explains the absence of signals in (in)direct detection experiments and colliders. A prediction of the model is that any future experiments of this kind will continue to produce null-results. In turn, a detection of a DM particle with mass below our predicted value would rule out our proposal of a heavy spin-2 as sole explanation for the observed DM. Alternative tests of our proposal could be based on its gravitational or its self-interacting nature. \paragraph{Outline of the paper.} Section~\ref{sec:BMreview} is dedicated to reviewing relevant details of bimetric theory. We provide its action, equations of motion and present the maximally symmetric solutions with the corresponding mass spectrum. In section \ref{sec:BMtoGR}, we discuss the two parameter regions for which the classical solutions of the theory resemble those of GR. The regime of parameters and energies where the bimetric action can be treated perturbatively is derived in the beginning of section~\ref{sec:heavymass}. Thereafter, we compute the cubic and quartic vertices in the spin-2 sector, verifying the absence of decay terms into massless gravitons. The phenomenology of spin-2 DM is explored in section~\ref{sec:pheno}. Finally, we discuss our results in section~\ref{sec:conclusions}. Additional supporting details can be found in the appendices.
\label{sec:conclusions} A heavy massive spin-2 field, whose gravitational interactions are described by the ghost-free bimetric theory, possesses all the desired features of a DM candidate. The extremely weak coupling of the spin-2 field to SM matter furthermore explains the absence of DM signals in dedicated detection experiments and collider searches. The Planck mass $\alpha m_g\sim \alpha\mpl$ of the second metric can range from $1$ to $10^4\,$TeV and the Fierz-Pauli mass for the spin-2 field is on the order of $\mfp\in[1,66]$~TeV. This narrow mass range for the DM candidate is one of the most distinct features of our model. Note however that the upper bound was obtained from the requirement of remaining in the perturbative framework. In principle, non-perturbative methods (which are presently unknown) could reveal that a larger spin-2 mass is also consistent with phenomenology. Another exceptional property is the universal decay of DM into all SM particles along with the absence of a decay channel into massless gravitons at tree level. Our scenario predicts that mass and interaction scale of the heavy spin-2 field are of the same order of magnitude and only slightly larger than the weak scale. The largeness of the physical Planck mass $\mpl$ is responsible both for suppressing the interactions of DM with baryonic matter and for bringing the theory close to GR. We have therefore not created any new hierarchies of energy scales and moreover related the puzzle of a large Planck scale to the extremely weak interactions of DM. The above constraints on bimetric parameters are consistent with all current gravity tests. They correspond precisely to the overlap of the regions $\mfp^2\gg\Lambda$ for the spin-2 mass and $\alpha\ll1$ for the ratio of Planck masses discussed in Sec.~\ref{sec:pointsol}. Indeed, the Compton wavelength of the spin-2 field is tiny, approximately between $10^{-19}~\text{cm}$ and $10^{-17}~\text{cm}$, resulting in typical Vainshtein radii of $r_V\sim 10^{-10}~\text{cm}$ for the Sun, and $r_V\sim 10^{-24}~\text{cm}$ for millimetre/sub-millimetre tests of the gravitational inverse-square law. These values ensure the validity of the linear approximation for all possible local gravity tests~\cite{Hoyle:2000cv, Yang:2012zzb}. Relative corrections to GR solutions (the ratio of the fifth force to the Newtonian force) involve a factor of $\alpha^2 \exp(-\mfp r)$, which is extremely small in our setup. The largest relative deviation from Newton's law, which would be accessible through the sub-millimetre tests, is thus at most $\sim 10^{-22} \exp(-10^{16})$, which is far beyond the reach of any experiment. In a very similar fashion, the smallness of the parameter $\alpha$ ensures that bimetric predictions for cosmology are essentially indistinguishable from the $\Lambda$CDM model and that perturbations remain stable. Bimetric theory in the above parameter region therefore resembles GR, but with an additional tensor field behaving like cold DM. The fact that static spherically symmetric solutions in the weak field approximation resemble GR solutions for $\alpha\gg 1$ and $\mfp^2\gg\Lambda$ may suggest that in this parameter range black holes do not have any specific features which would distinguish them from those of GR. For instance, the instability found for spherically symmetric black holes~\cite{Babichev:2013una,Brito:2013wya} --- which is clearly a distinctive feature of bimetric theory, since in GR such black holes are stable --- disappears for large $\mfp$, since the instability range is limited to $\mfp\lesssim r_S^{-1}$. On the other hand, the absence of Birkhoff's theorem in bimetric theory allows for the existence of hairy black holes, see~\cite{Volkov:2012wp,Brito:2013xaa} for particular examples and~\cite{Babichev:2015xha} for a recent discussion. It is therefore feasible that hairy black holes are present in our scenario, possibly giving rise to distinct observational features of bimetric theory. It remains to be answered how different these hairy black holes can be from GR black holes, but this question lies beyond the scope of our paper. Another distinct property of spin-2 DM are its enhanced self-interaction terms whose form is fixed by the ghost-free structure of the bimetric potential. Self-interacting DM is known to produce observable effects in collisions of galaxy clusters but so far the constraints are not very stringent~\cite{Harvey:2015hha}. Moreover, DM self-interactions would induce distortions in the DM power spectrum at small scales while leaving the baryonic one untouched. However, these interactions are mediated by $\dM$ itself, which confines the effectiveness of the corresponding force to risible length scales in astronomical or cosmological terms. Thus, current constraints on the self-interaction cross section are of little relevance in our case, despite the strength of the $1/\alpha$ enhancement, since in practice the DM particles never feel each other. Let us finally point out that, in addition to the aforementioned characteristics of our model, the decay of spin-2 DM would in principle exhibit a slightly different spectrum of secondaries (especially neutrinos), compared to the decay of, for instance, a scalar singlet. Indeed, since the $\delta M_{\mu\nu}$ field couples directly to the energy-momentum tensor of the SM, the two-vector final state (such as $ZZ$ or $W^+W^-$) will be mostly transversal, that is, will carry spin 2. This is in contrast to DM being a scalar, since in that case only longitudinal states can be produced (at tree level, but other polarisations can appear through higher dimensional operators). It can also differ from models with a Kaluza-Klein massive graviton, since in those constructions the transversal-to-longitudinal ratio can be different (depending on the structure of the extra dimensional setup), see for example~\cite{Han:2015cty}. However, this is a very model-dependent statement and we emphasise that in bimetric theory all predictions are fixed by demanding consistency of the theory. Once the vector bosons decay, the final spectra of secondary neutrinos will bear information about the spin in the guise of peculiar spectral features, see for instance the discussion in~\cite{Garcia-Cely:2016pse}. These features, however, are very small and hardly within the reach of current experiments. \vspace{0.5cm}
16
7
1607.03497
We provide further details on a recent proposal addressing the nature of the dark sectors in cosmology and demonstrate that all current observations related to Dark Matter can be explained by the presence of a heavy spin-2 particle. Massive spin-2 fields and their gravitational interactions are uniquely described by ghost-free bimetric theory, which is a minimal and natural extension of General Relativity. In this setup, the largeness of the physical Planck mass is naturally related to extremely weak couplings of the heavy spin-2 field to baryonic matter and therefore explains the absence of signals in experiments dedicated to Dark Matter searches. It also ensures the phenomenological viability of our model as we confirm by comparing it with cosmological and local tests of gravity. At the same time, the spin-2 field possesses standard gravitational interactions and it decays universally into all Standard Model fields but not into massless gravitons. Matching the measured DM abundance together with the requirement of stability constrains the spin-2 mass to be in the 1 to 100 TeV range.
false
[ "Massive spin-2 fields", "Dark Matter searches", "Dark Matter", "General Relativity", "standard gravitational interactions", "massless gravitons", "the heavy spin-2 field", "baryonic matter", "gravity", "Standard Model", "a heavy spin-2 particle", "the spin-2 field", "experiments", "all Standard Model fields", "signals", "the spin-2 mass", "ghost-free bimetric theory", "cosmology", "TeV", "their gravitational interactions" ]
8.623079
-1.507675
54
12406023
[ "De Paolis, F.", "Gurzadyan, V. G.", "Nucita, A. A.", "Chemin, L.", "Qadir, A.", "Kashin, A. L.", "Khachatryan, H. G.", "Sargsyan, S.", "Yegorian, G.", "Ingrosso, G.", "Jetzer, Ph.", "Vetrugno, D." ]
2016A&A...593A..57D
[ "Triangulum galaxy viewed by Planck" ]
12
[ "Dipartimento di Matematica e Fisica \"Ennio De Giorgi\", Università del Salento, via per Arnesano, 73100, Lecce, Italy ; INFN, Sezione di Lecce, via per Arnesano, 73100, Lecce, Italy", "Center for Cosmology and Astrophysics, Alikhanian National Laboratory and Yerevan State University, Yerevan, Armenia; SIA, Sapienza University of Rome, 00185, Rome, Italy", "Dipartimento di Matematica e Fisica \"Ennio De Giorgi\", Università del Salento, via per Arnesano, 73100, Lecce, Italy; INFN, Sezione di Lecce, via per Arnesano, 73100, Lecce, Italy", "LAB, CNRS UMR 5804, Université de Bordeaux, 33270, Floirac, France", "School of Natural Sciences, National University of Sciences and Technology, 44000, Islamabad, Pakistan", "Center for Cosmology and Astrophysics, Alikhanian National Laboratory and Yerevan State University, Yerevan, Armenia", "Center for Cosmology and Astrophysics, Alikhanian National Laboratory and Yerevan State University, Yerevan, Armenia", "Center for Cosmology and Astrophysics, Alikhanian National Laboratory and Yerevan State University, Yerevan, Armenia", "Center for Cosmology and Astrophysics, Alikhanian National Laboratory and Yerevan State University, Yerevan, Armenia", "Dipartimento di Matematica e Fisica \"Ennio De Giorgi\", Università del Salento, via per Arnesano, 73100, Lecce, Italy; INFN, Sezione di Lecce, via per Arnesano, 73100, Lecce, Italy", "Physik-Institut, Universität Zürich, Winterthurerstrasse 190, 8057, Zürich, Switzerland", "Department of Physics, University of Trento; and TIFPA/INFN, 38123 Povo, Trento, Italy" ]
[ "2018A&A...609A.131G", "2019A&A...629A..87D", "2019IJMPD..2830021B", "2019IJMPD..2840016A", "2019IJMPD..2850088T", "2019MPLA...3450308A", "2019PhRvD.100d3028Q", "2020PhRvD.101h3016Z", "2021ApJS..253...53W", "2021EPJC...81..827T", "2022A&A...664A..30T", "2023Symm...15..160T" ]
[ "astronomy" ]
4
[ "galaxies: general", "galaxies: individual: M 33", "galaxies: halos", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1051/0004-6361/201628780", "10.48550/arXiv.1607.08099" ]
1607
1607.08099_arXiv.txt
A temperature asymmetry has been detected in {\it Planck} data toward the M31, Cen A, and M82 galaxies, always aligned with respect to the expected galaxy spin (\citealt{depaolis2011,depaolis2014,depaolis2015,gurzadyan2015}). The aim of this paper is to investigate whether the same phenomenon also occurs for M33 (also known as the Triangulum galaxy), which is an appropriate object on which to test our method, and on which we can trace the dynamics of the baryonic galactic halo (which at large scales also has to reflect the dark matter contribution) in a model-independent way because it is sufficiently extended and useful multi-wavelength observations are available. M33 (or NGC 598) is the nearest late-type spiral galaxy that lies at a distance of only about 840 kpc \citep{magrini2007,gieren2013}. It is the third largest member (after M31 and the Milky Way) of the Local Group (see \citealt{rg} for the dynamics and substructure of the Local Group), and the coordinates of its center are RA $01^h 33^m 50.9^s$ and Dec $30\degr 39\arcmin 37 \arcsec$. It is classified as an SAcd galaxy, that is, a late-type spiral with a weak bar (\citealt{corbelliwalterbos2007}), no clear evidence of any bulge component (see, e.g., \citealt{bothan1992}), and relatively loosely wound arms \citep{buta2007}. At the M33 distance $1\arcmin$ corresponds to 244 pc ($1\degr\simeq 14.67$ kpc), which allows us to study this galaxy with a great degree of accuracy. The relatively small inclination angle $i\simeq 56\degr$ \citep{vandenbergh2000} allows us to obtain a comprehensive view of the M33 galaxy, which is a key advantage, for example, for studying the correlation of the velocity field with the galaxy geometry. For a recent and comprehensive review about M33 we refer to \cite{hodge2012}. \begin{figure}[h!] \centering \includegraphics[width=0.46\textwidth]{Fig11degree.jpg} \caption{{\it Planck} field toward the M33 galaxy in the 70 GHz band. The pixel color gives the temperature excess in $\mu$K with respect to the mean CMB temperature. The optical extension of the M33 galaxy is indicated by the inner ellipse with major and minor axes of $60\arcmin$ and $35\arcmin$, respectively. The four quadrants A1, A2, A3, and A4 are used in our analysis.} \label{fig1} \end{figure} The M33 disk is rotating with a maximum circular velocity of about $120-130$ km s$^{-1}$ and has a rising profile (see, e.g., \citealt{corbellisalucci2000,corbelli2014,kam2015}), indicating that the outer regions of the galaxy possess substantial mass. The observations at wavelength of 21 cm show the detailed spatial and kinematic structure of the neutral hydrogen in M33, which extends farther out than the stellar component. Interestingly, at least $18\%$ of the HI gas is found beyond the star-forming disk. There is some evidence of a faint halo component (\citealt{mcconnachie2010,cockcroft2013}) that was previously undetected (see, e.g., \citealt{trinchieri1988,barker2007}). The study in the microwaves of the M33 disk and halo within $4\degr$ is the aim of the following sections, particularly with the aim of investigating their rotation.
Similar to other nearby galaxies analyzed in some previous papers, M33 also shows a temperature asymmetry in {\it Planck} data with respect to the M33 minor axis projected onto the sky plane. This effect is clearly visible in all considered {\it Planck} bands at 70, 100, and 143 GHz, as well as in the best foreground-corrected SMICA map. As shown by the good correlation with the HI velocity field within about $0.5\degr$, this is likely due to gas or dust emission in the M33 disk, modulated by its rotation. After peaking in the 70 GHz band at about $0.5\degr$, the microwave signal decreases between $0.5\degr$ and $1\degr$, but then it increases again and peaks in the annulus between $2\degr$ and $3\degr$ (see Fig. \ref{fig6}). The detected temperature asymmetry is much more extended than the visible M33 disk and seems to be present up to a galactocentric distance of about $3\degr$. This apparently odd behavior indicates either a change in the emission geometry or a drastic change of the dominant emission mechanism, even if the situation might be complicated by contaminating effects, such as gas flowing into the M33 galaxy, and/or related to a past M33-M31 interaction (see, in particular, \citealt{bekki2008,wolf2013}). Five possible emission mechanisms might explain the observed temperature asymmetry: ($i$) free-free emission, ($ii$) synchrotron emission, ($iii$) anomalous microwave emission (AME) from dust grains, ($iv$) Sunyaev-Zel'dovich (SZ) effect, in particular the kinetic-SZ effect, and ($v$) cold gas clouds populating the halo of M33. The absence of a substantial dependence of the observed temperature asymmetry on the considered {\it Planck} bands together with the fact that the hot and cold side are toward the expected directions indicates that it has to be modulated by the Doppler effect induced by the M33 spin. \begin{figure}[h!] \centering \includegraphics[width=0.48\textwidth]{M33halotilted.pdf} \caption{Temperature asymmetry in $\mu$K, in the 70 GHz {\it Planck} band, of region A3 and A4 with respect to A1 and A2 for the second, third, and fourth annuli around M33 with PA$=-13\degr$.} \label{fig7} \end{figure} The dominating emission mechanism in the microwave regime is not clear as yet. Mechanisms ($i$), ($ii$), and ($iv$) require the presence of a hot plasma with cosmic-ray electrons (CRE) around M33, however, and give rise to an emission spectrum with a rather steep dependence on the frequency (see, e.g., \citealt{bennett2003,planck2014}). This plasma, expected to emit mainly in X-rays, is hard to detect since the emission measure scales with the square of the electron density, and when the galaxy surface brightness drops to the level of the X-ray background, mainly due to the Milky Way, further detection is impossible. A hot gaseous halo (or corona) has been detected only for a few spiral galaxies. In our Galaxy, \cite{gupta} probed the warm-hot phase (at a temperature $\simeq 10^6$ K) of the Galactic halo through the Chandra detection of $O_{VII}$ and $O_{VIII}$ absorption lines (with redshift $z=0$) toward known AGNs and inferred a warm halo extending up to $\simeq 100$ kpc and possibly containing about $10^{10}~M_{\odot}$ in gaseous form \citep{miller2013}. This gaseous halo probably rotates with a velocity lower than the Local Standard of Rest velocity $v_{\rm LSR}$ \citep{Hodges-Kluck2016}. For UGC12591, a hot halo has been detected that extends up to 110 kpc from the galactic center (\citealt{dai2012}). For M33, although a hot-diffuse gas component is clearly visible in the XMM-Newton EPIC data up to about $30\arcmin$ (\citealt{pietsch2003}), it cannot be excluded that it extends farther out and, moreover, investigation of the CRE propagation length shows that its value is even higher than that in M31 (\citealt{berkhuijsen2013}). AME, option ($iii$), has been observed in various interstellar environments such as in the ISM \citep{miville2008} and in dark clouds \citep{watson2005}. Several processes can excite dust grain emission, provided they are present in galactic halo environments (see below): IR and radio photons, gas or dust interaction, formation of $H_2$ molecules on the dust grain surface, photoelectric emission (see, e.g., \citealt{draine1998}). Option ($v$) cannot be excluded at present because cold gas clouds (with or without a dust component) may populate the M33 halo, giving rise, if it were to rotate, to a certain temperature asymmetry through Eq. (\ref{eq1}). These clouds may be optically thin or thick to their own sub-millimeter radiation, depending on their mass and temperature. When the clouds are in virial equilibrium, the cloud optical thickness may be estimated through the relation \begin{equation} \langle\tau\rangle\simeq 6\times 10^{-2} \left(\frac{T_c}{2.7~ {\rm K}} \right)^2 \left(\frac{M_c}{M_J} \right)^{-1}, \end{equation} where $M_c$ and $T_c$ are the cloud mass and temperature, and $M_J$ is the mass of Jupiter. Large clouds are therefore expected to be optically thin, while only extremely low massive clouds (with $M_c\leq 0.1 M_J$) may be optically thick. The expected effect in the microwave bands is therefore very small. Cold clouds, such as those implied by BOOMERANG \citep{boomerang2010}, {\it Planck} \citep{planck2011}, and {\it Herschel} \citep{juvela2010,juvela2012}, also called{\it Herschel} cold clouds (HCC) and studied mainly toward the Small Magellanic Clouds (see \citealt{nieuwenhuizen2012} and references therein), might also be present in the M33 halo and give rise to a temperature asymmetry in the {\it Planck} bands. We also mention that HI clouds have been discovered in the outer regions of M33 \citep{grossi2008,keenan2016}, and it has been suggested that they play a role in feeding the star formation in the disk of M33. The observed gaseous features extend up to 38 kpc from the galaxy center, and the HI mass beyond the M33 disk contains about $18\%$ of the estimated total HI mass \citep{putman2009}. Until now, {\it Herschel} data of M33, within the project HERM33ES ({\it Herschel} M33 Extended Survey), have been analyzed and reported in the literature only up to a galactocentric distance of about $0.5-0.6\degr$ and show a cold dust component with temperature $\simeq 14$ K between about 2 kpc and 6 kpc, and both SPIRE and PACS maps show a weak emission extending outside of 8 kpc \citep{kramer2010}. For the extension of the microwave emission toward M33 and its connection with some galactic population, we mention that in addition to the possible hot corona of CRE and plasma, there is another component that has recently been discussed in the literature: the RGB star halo population of M33, which was discovered by using the Pan-Andromeda Archeological Survey (PAndAS). This is found to have an exponential scale length of about $20$ kpc (about $1.4\degr$), as obtained by \cite{cockcroft2013}, and might account for the emission in the microwave bands through the mechanisms discussed above. Even though RGB stars cannot generate an asymmetry in the microwave temperature, a gas or dust component (as that described above) might be present in the outer region of the M33 galaxy, in association with the RGB stars. Galactic halos do contain the fossil record of galaxy assembly, and clarifying their composition is extremely important to understand the processes that led to the formation and subsequent evolution of these structures. The degree to which galactic halos rotate with respect to the disks is also a relevant and extremely difficult question to answer (see, e.g., \citealt{courteau2011,deason2011}). All this indicates the need for more detailed, extended, and dedicated surveys.
16
7
1607.08099
We used Planck data to study the M 33 galaxy and find a substantial temperature asymmetry with respect to its minor axis projected onto the sky plane. This temperature asymmetry correlates well with the HI velocity field at 21 cm, at least within a galactocentric distance of 0.5°, and it is found to extend up to about 3° from the galaxy center. We conclude that the revealed effect, that is, the temperature asymmetry and its extension, implies that we detected the differential rotation of the M 33 galaxy and of its extended baryonic halo.
false
[ "its extended baryonic halo", "a substantial temperature asymmetry", "the galaxy center", "respect", "the sky plane", "This temperature asymmetry", "the temperature asymmetry", "Planck data", "the HI velocity field", "a galactocentric distance", "the M 33 galaxy", "the differential rotation", "its minor axis", "the revealed effect", "21 cm", "its extension", "0.5", "We", "about 3", "it" ]
13.527616
2.600409
-1
12406558
[ "Auddy, Sayantan", "Basu, Shantanu", "Valluri, S. R." ]
2016AdAst2016E..13A
[ "Analytic Models of Brown Dwarfs and the Substellar Mass Limit" ]
27
[ "-", "-", "-" ]
[ "2017ApJ...846...97G", "2017RAA....17..122A", "2018MNRAS.478.2113H", "2019ApJ...871..227F", "2019EPJC...79.1030R", "2020ApJ...888..102L", "2020PhR...876....1O", "2021PhRvD.103f4032B", "2021PhRvD.104j4058W", "2021PhRvL.126p1101L", "2022ApJ...928...91C", "2022ApJ...929..117C", "2022Univ....8..647K", "2022ZNatA..77..515N", "2023CQGra..40s5021P", "2023EPJC...83..492G", "2023JCAP...10..057L", "2023MNRAS.518.1484Y", "2023PhRvD.107d3012B", "2023PhRvD.108b4016K", "2023ZNatA..78..325N", "2024IJMPA..3950026B", "2024JCAP...04..038A", "2024JCAP...04..082I", "2024MNRAS.52710737F", "2024PhRvD.109j4049C", "2024Univ...10..217W" ]
[ "astronomy" ]
11
[ "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1155/2016/5743272", "10.48550/arXiv.1607.04338" ]
1607
1607.04338_arXiv.txt
S_1 = \frac{k_B N_0}{\mu_{1\rm{mod}}} \ln \frac{T^\frac{3}{2}}{\rho_1} + C_1, \end{equation} where \begin{equation} \frac{1}{\mu_{1\rm{mod}}}=\left(\frac{1}{\mu_1}+\frac{3}{2}\frac{x_{H^{+}}(1-x_{H^{+}})}{2-x_{H^{+}}}\right), \end{equation} and $\mu_1$ is different for each model in this region and is calculated using $x_{H^{+}}$ from Table 1. However, the contributions to entropy due to radiation and the degenerate electrons (see Eq 2-145 in \cite{Clay68}) are negligible in the range of temperature and density applicable for brown dwarfs. Based on a similar argument, the analytic expression for the entropy of non-ionized molecular hydrogen and helium mixture at the photosphere is expressed as \begin{equation}\label{the exterior entropy} S_2 = \frac{k_B N_0}{\mu_2} \ln \frac{T^\frac{5}{2}}{\rho_2} + C_2, \end{equation} where $\frac{1}{\mu_2} = \frac{X}{2}+\frac{Y}{4}$ is the mean molecular weight for the hydrogen and helium mixture. The expression for entropy $S_2$ is derived using the first law of thermodynamics (Appendix B) and the relation of the internal energy of diatomic molecules ($ U=\frac{5}{2} \frac{k_BN_0T}{\mu_2}$). Here we have considered only five degrees of freedom as the temperatures are just sufficient to excite the rotational degrees of H$_2$ but the vibrational degrees remain dormant. It should be noted that Eqs. (\ref{the interior entropy}) and (\ref{the exterior entropy}) are just simplified forms of the entropy expressions presented in \cite{Burs93} and \cite{Ste91}. Thus the entropy in the two phases is dominated by contributions from the ionic and molecular gas, respectively. Using the same argument as \cite{Burs93}, that the two regions of different temperature and density are separated by a phase transition of order one we can estimate the surface temperature. Using the expression for the degeneracy parameter $\psi$ from Eq. (\ref{the degeneracy parameter}) we can simplify Eq. (\ref{the interior entropy}) to be \begin{equation}\label{entropy1mod} S_1 = \frac{3}{2}\frac{k_B N_0}{\mu_{1\rm{mod}}}( \ln \psi +12.7065) + C_1. \end{equation} Furthermore, the jump of entropy \begin{equation}\label{entropyjump} \Delta \sigma=\frac{S_2-S_1}{k_BN_0}, \end{equation} (see Table 1) for the phase transition at each point of the coexistence curve of PPT \citep{Saumon89}, is used to estimate the relation $\mid C_1-C_2\mid$ between the two constants in Eqs. (\ref{the interior entropy}) and (\ref{the exterior entropy}). For $T=T_{\rm{eff}}$ and $\rho_2=\rho_e$ in Eq (\ref{the exterior entropy}), we can use Eqs. (\ref{entropy1mod}) and (\ref{the exterior entropy}) in Eq. (\ref{entropyjump}) and the value of $\mid C_1-C_2\mid$ to obtain a wide range of possible values of surface temperature $T_{\rm{eff}}$ in terms of the degeneracy parameter and photospheric density $\rho_e$: \begin{equation}\label{effective temperature} T_{\rm{eff}}= b_1\times 10^{6} \rho_e^{0.4} \psi^v \:\:\rm{K}. \end{equation} The values of the parameters $b_1$ and $v$ for different models are shown in Table \ref{tbl:model}. According to the Chabrier model, the critical temperature $ \sim 1.53 \times 10^4$ K and critical density $ \sim 0.35 $ g cm$^{-3}$ marks the end of the phase transition with $ \Delta \sigma=\frac{\Delta S}{k_BN}=0 $. In the following discussion we briefly summarise the steps from \cite{Burs93}. We replace Eq. (2.50) in \cite{Burs93} by Eq. (\ref{effective temperature}) to estimate the surface luminosity. As an example we select a particular phase transition point (model D) and show the derivation of surface luminosity using hydrostatic equilibrium and the ideal gas approximation. The photosphere of a brown dwarf is located at approximately the $\tau = \frac{2}{3}$ surface, where \begin{equation}\label{Tau} \tau=\int_r^{\infty} \kappa_R \rho \textit{ dr}. \end{equation} is the optical depth. Using the general equation for hydrostatic equilibrium, $dP=-(GM/r^2)\rho \, \textit{dr}$, and Eq. (\ref{Tau}), the photospheric pressure can be expressed as \begin{equation}\label{external pressure 1} P_e=\frac{2}{3}\frac{GM}{\kappa_R R^2}, \end{equation} where $\kappa_R $ is the Rosseland mean opacity and the other variables have their standard meanings. Furthermore, our EOS (Eq. (\ref{the power law eq of state})) in the approximation of negligible degeneracy pressure near the photosphere gives the photoshperic pressure as \begin{equation}\label{external pressure 5} P_e=\frac{\rho_e N_Ak_B T_{\rm{eff}}}{\mu_2}. \end{equation} Now using the expression for radius $R$ (Eq. \ref{radius}) in Eq. (\ref{external pressure 1}), we can calculate the external pressure $P_e$ as a function of $M$ and $\psi$: \begin{equation}\label{external pressure} P_e= \frac{11.2193 \,\rm bar}{\kappa_R}\left(\frac{M}{M_{\odot}}\right)^{5/3} \frac{\mu_e^{10/3}}{(1+\gamma +\alpha\psi)^2} . \end{equation} On using Eq. (\ref{external pressure}) in Eq. (\ref{external pressure 5}) and substituting $T_{\rm{eff}}$ for model D with $b_1=2.00 $ and $v=1.60$ from Table \ref{tbl:model}, the effective density $\rho_e$ can be expressed as a function of $M$ and $\psi$: \begin{equation}\label{the effective density} \rho_e^{1.40}=\frac{6.89811}{\kappa_R N_A k_B} \left(\frac{M}{M_{\odot}}\right)^{5/3} \frac{\mu_e^{10/3}\mu_2 \; \rm {g/cm^3} }{(1+\gamma+\alpha\psi)^{2}\psi^{1.58}}. \end{equation} Substituting the expression for $\rho_e$ from Eq. (\ref{the effective density}) in Eq. (\ref{effective temperature}) we derive the expression for effective temperature for model D as a function of $M$ and $\psi$: \begin{equation}\label{t effective} T_{\rm{eff}}=\frac{2.57881 \times 10^4 \, \rm K }{\kappa_R^{.2856}} \left(\frac{M}{M_{\odot}}\right)^{0.4764} \frac{\psi^{1.1456}}{(1+\gamma+\alpha\psi)^{0.5712}} . \end{equation} Similarly, the surface temperature can be evaluated for all the other models. Since the procedure is same for all the models in Table \ref{tbl:model}, we just show one calculation. For this range of surface temperatures, the Stefan-Boltzmann law, $L=4\pi R^2 \sigma T_{\rm{eff}}^4$, yields a set of possible values of the surface luminosity $L$ as a function of the degeneracy parameter $\psi$. The luminosity for model D using Eq. (\ref{radius}) and Eq. (\ref{t effective}) is \begin{equation}\label{luminosity model D} L=\frac{0.41470\times L_{\odot}}{\kappa_R ^{1.1424}} \left(\frac{M}{M_{\odot}}\right)^{1.239} \frac{\psi^{4.5797}}{(1+\gamma+\alpha\psi)^{0.2848}}, \end{equation} where $\sigma$ is the Stefan-Boltzmann constant. Substellar objects below the main sequence mass gradually evolve towards complete degeneracy and a state of stable equilibrium as their luminosity decreases over time. In the following sections we show that the degeneracy parameter $\psi$ is a function of time and it evolves toward $\psi =0$ over the lifetime of brown dwarfs. This gives us an estimate of the luminosity at different epochs of time. \begin{deluxetable}{lcccccccc} \tabletypesize{\footnotesize} \tablewidth{-4pt} \tablecaption{Effective temperature for different phase transition points\tablenotemark{a} \label{tbl:model}} \tablehead{ \colhead{${\rm Model} $} & \colhead{$\log T$ (K)} & \colhead{$P$(Mbar)} & \colhead{$\rho_1$ (g cm$^{-3}$)} & \colhead{$\rho_2$ (g cm$^{-3}$)} & \colhead{$\Delta \sigma$ } & \colhead{$2x_{H^{+}}$ } & \colhead{$b_1 $} & \colhead{$v$} } \startdata A &3.70 & 2.14 & 0.75 & 0.92 & 0.62 & 0.48 & 2.87 & 1.58 \\ B &3.78 & 1.95 & 0.70 & 0.88 & 0.59 & 0.50 &2.70 & 1.59 \\ C &3.86 & 1.62 & 0.64 & 0.80 & 0.54 & 0.50 & 2.26 & 1.59 \\ D &3.94 & 1.39 & 0.58 & 0.74 & 0.51 & 0.51 & 2.00 & 1.60 \\ E &4.02 & 1.13 & 0.51 & 0.65 & 0.46 & 0.52 &1.68 & 1.61 \\ F &4.10 & 0.895 & 0.43 & 0.55 & 0.42& 0.50& 1.29 & 1.59 \\ G &4.18 & 0.631 & 0.35 & 0.38 & 0.14& 0.33&0.60 & 1.44 \\ H &4.185 & 0.614 & 0.35 & 0.35 & 0.00& 0.18& 0.40 & 1.30 \\ \enddata \tablenotetext{a}{The phase transition points are taken from \citet{Chab92}. This gives the possible range of surface temperature depending on the phase transition points. For different values of temperature and density at which the phase transition takes place, the effective surface temperature is calculated using Eq. (\ref{effective temperature}).} \end{deluxetable} \subsection{Validity of PPT in brown dwarfs} There is a distinction between the temperature-driven PPT with a critical point at $ \sim 0.5$ Mbar and between $10000$ K and $20000$ K as predicted by the chemical models \citep{Saumon95}, and the pressure-driven transition from an insulating molecular liquid to a metallic liquid with a critical point below 2000 K at pressures between $1$ and $3$ Mbars. The latter is predicted e.g., by \cite{Lor09}, \cite{Maz10}, and \cite{Mor10} based on the ab initio simulations . \cite{Lor09} rule out the presence of PPT above $10000$ K and give an estimate of the critical points for the transition at $T_c = (1400 \pm 100)\, \rm K$, $P_c = 1.32 \pm 0.1$ Mbar, $\rho_c = 0.79 \pm 0.05$ $\rm g/cm^3$. Similarly \cite{Mor10} estimated the critical point of the transition at a temperature near $2000$ K and pressure near $1.2$ Mbar. Signatures of pressure-driven PPT in a cold regime below $2000$ K are obtained by \cite{Knu15}. Figure 1 in \cite{Knu15} shows the melting line (black) as well as the different predictions for the coexistence lines for the first order transition (green curves). Brown dwarf interior temperatures are far above these estimates for a first order transition from the insulating to the metallic system. The same is true for Jupiter. Of course a continuous transition may be possible in Jupiter and brown dwarfs, but a first order transition may not be possible. Thus the determination of the range of temperature of this transition provides a much needed benchmark for the theory of the standard models for the internal structures of the gas-giant planets and low mass stars.
This paper presents a simple analytic model of substellar objects. A focus of the paper was to revisit both the development and the shortcomings of the theoretical understanding of the physics governing the evolution of low mass stars and substellar objects over the last 50 years. We have also made some modifications to the existing models to better explain the physics using analytic forms. Although observational constraints hinder our understanding, a simple analytic model can answer many questions. We have summarized the method of determining the minimum mass for sustained hydrogen burning. Objects in the mass range $0.064M_\odot-0.087M_\odot$ mark this critical boundary between brown dwarfs and the main sequence stars. \\ We have derived a general equation of state using polylogarithm functions \citep{Val10} to obtain the $P-\rho$ relation in the interior of brown dwarfs. The inclusion of the finite temperature correction gives us a much more complete and sophisticated analytic expression of the Fermi pressure (Eq. \ref{pressure eqn3}). The application of this relation can extend to other branches of physics, especially for semiconductor and thermoelectric materials \citep{Molli2011}.\\ The estimate of the surface luminosity is a challenge given our limitations in understanding the physics inside such low mass objects. Also, it is still an open question if a phase transition actually occurs in a brown dwarf. The results of modern day simulations \citep{Yan15, Mor10} do raise doubts about the relevance of phase transitions in the brown dwarf scenario. We are not aware of well defined analytic models that have a unique way of estimating the surface luminosity apart from using the PPT technique as given in \cite{Burs93}. In this work, rather than considering a single value for the phase transition point we have used the entire range of temperatures from the phase transition coexistence table \citep{Chab92}. These are within the uncertainty range of the critical temperature of the PPT as proposed by the recent simulations \citep{Yan15}. Thus, considering the large uncertainties involved in such models, this range of values of the minimum mass is much more acceptable than a single distinct transition mass. However, the next step forward is to develop an analytic model for surface temperature that is independent of the PPT. % We estimate that $\simeq 5 \% $ of stars take more than $10^8$ yr to reach the main sequence, and $\simeq 11 \% $ of stars take more than $10^7$ yr to reach the main sequence (Table 2). The stars in these categories have mass very close to the minimum hydrogen burning limit, and will eventually settle into an extremely low luminosity main sequence with $L/L_{\odot}$ in the range $\approx 10^{-5} - 10^{-4}$. The very low luminosity non-main-sequence hydrogen burning in substellar objects and the pre-main-sequence nuclear burning in very low mass stars are very interesting to study further, and our simplified model can certainly be improved in its ability to estimate the time evolution.
16
7
1607.04338
We present the current status of the analytic theory of brown dwarf evolution and the lower mass limit of the hydrogen burning main sequence stars. In the spirit of a simplified analytic theory we also introduce some modifications to the existing models. We give an exact expression for the pressure of an ideal non-relativistic Fermi gas at a finite temperature, therefore allowing for non-zero values of the degeneracy parameter ($\psi = \frac{kT}{\mu_{F}}$, where $\mu_{F}$ is the Fermi energy). We review the derivation of surface luminosity using an entropy matching condition and the first-order phase transition between the molecular hydrogen in the outer envelope and the partially-ionized hydrogen in the inner region. We also discuss the results of modern simulations of the plasma phase transition, which illustrate the uncertainties in determining its critical temperature. Based on the existing models and with some simple modification we find the maximum mass for a brown dwarf to be in the range $0.064M_\odot-0.087M_\odot$. An analytic formula for the luminosity evolution allows us to estimate the time period of the non-steady state (i.e., non-main sequence) nuclear burning for substellar objects. Standard models also predict that stars that are just above the substellar mass limit can reach an extremely low luminosity main sequence after at least a few million years of evolution, and sometimes much longer. We estimate that $\simeq 11 \%$ of stars take longer than $10^7$ yr to reach the main-sequence, and $\simeq 5 \%$ of stars take longer than $10^8$ yr.
false
[ "brown dwarf evolution", "an ideal non-relativistic Fermi gas", "evolution", "substellar objects", "the non-steady state", "stars", "non-zero values", "Fermi", "surface luminosity", "an extremely low luminosity main sequence", "(i.e., non-main sequence", "the substellar mass limit", "the lower mass limit", "the luminosity evolution", "\\psi", "Standard models", "a finite temperature", "its critical temperature", "the Fermi energy", "the main-sequence" ]
7.734277
12.648236
-1
12464732
[ "Herrera, Ramón", "Hipólito-Ricaldi, W. S.", "Videla, Nelson" ]
2016JCAP...08..065H
[ "Instability in interacting dark sector: an appropriate holographic Ricci dark energy model" ]
13
[ "Instituto de Física, Pontificia Universidad Católica de Valparaíso, Avenida Brasil 2950, Casilla 4059, Valparaíso, Chile", "Departamento de Ciências Naturais, Universidade Federal do Espírito Santo, Rodovia BR 101 Norte, km. 60, São Mateus, Espírito Santo, Brasil ; Department of Physics, McGill University, Montréal, QC, H3A 2T8, Canada;", "Departamento de Física, Universidad de Chile, FCFM, Blanco Encalada 2008, Santiago, Chile" ]
[ "2016JCAP...08..072F", "2017PhRvD..95l3521S", "2017arXiv170204244M", "2018BrJPh..48..364L", "2018PhRvD..98d3517Y", "2019EPJC...79..889L", "2020JPhCS1558a2004D", "2021EPJC...81...31C", "2021PhRvD.103h3511L", "2021PhRvD.104j3508L", "2022PhRvD.106d3527L", "2023EL....14269001Z", "2024EPJC...84..558C" ]
[ "astronomy" ]
4
[ "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1088/1475-7516/2016/08/065", "10.48550/arXiv.1607.01806" ]
1607
1607.01806_arXiv.txt
It is well known that the measurements of the luminosity redshift of supernovae (type Ia) have proportioned growing evidence for a phase of accelerated expansion of current universe\cite{C1,C2}. However, other evidences of this accelerated expansion come from baryon acoustic oscillations \cite{C3}, anisotropies of the cosmic microwave background (CMB)\cite{C4}, and among other \cite{C5}, have confirmed this scenario. In order to obtain a phase of accelerated expansion in Einstein's General Relativity it is necessary that the cosmological background dynamics be dominated by some exotic component with a negative pressure, known as dark energy (DE). Assuming that DE contributes to an important fraction of the content of the observable universe, it is instinctive and natural from the field theory to assume its interactions with other fields, e.g., dark matter (DM). It is well known that a suitable interaction between DE and DM can provide an novel mechanism to alleviate the cosmic coincidence problem \cite{C6,C7}. On the other hand, these models affect the structure formation and hence provide a different way to change the predictions of non-interacting models. Regarding this point, the interaction between DE and DM have been studied considering different types of observational data sets, see Refs.\cite{C8,C9}. For more comprehensive references of models with interacting DE and DM, see Refs.\cite{I1,I2}. The study of structure formation in models of DE and DM, through the cosmological perturbations theory, plays a fundamental role when these models are confronted with the observations \cite{AC1}. These models imprint a signature on the CMB power spectrum \cite{AC2,AC3} and also the space of parameters is modified\cite{AC3,AC4} . For this reason the analysis of the cosmological perturbations is important and also need to be well-behaved. In particular for interacting models, the background dynamics with adiabatic initial conditions and the perturbation theory were analyzed in Ref.\cite{C10}. Here, the perturbative dynamics realizes unstable growing modes. A further analysis in models with an interacting DE component together with a constant equation of state (EoS) $w$, was considered in Ref. \cite{C11}. Here the authors found that perturbations were unstable and with a rapid growth of DE fluctuations. To avoid a possible conflict with the perturbative dynamics when the EoS parameter $w$ crosses the value $w=-1$, in Ref.\cite{Kunz:2006wc}, the authors considered a new variable associated to the divergence of the velocity field. However, the new perturbations equations have a term associated to the pressure perturbation and, therefore, an adiabatic speed of sound. In this way, the authors introduced a free parameter in the adiabatic speed of sound, avoiding the divergences in the perturbations. On the other hand, considering that observational-data tests of the $\Lambda$ cold dark matter ($\Lambda$CDM) model are not accurate enough to rule out adequately the large diversity of alternative DE models in the literature that have been proposed to account for the data. In modern cosmology, one can test the $\Lambda$CDM to describe an adequate DE model, assuming the DE as an effective fluid (or a scalar field) and considering its EoS as a free and dynamics parameter. Some candidates for this DE are the holographic models (HM) which give a specific classes of dynamic approaches to solve the cosmic coincidence problem and is another alternative to the standard $\Lambda$CDM model. These models are motivated from the Holographic principle which has its origin in the black hole and string theories \cite{HP1}. These HM have a direct connection between an ultraviolet and infrared cutoff \cite{HM1,HMM1,HMM2}. In this context, the infrared cutoff corresponds to a cosmological length scale, and by the other hand this connection between cutoffs ensures that the energy density does not exceed the energy in a given volume of a black hole of the equal size. Regarding the HM with interaction between DE and DM, this kind of models were studied in Refs.\cite{H0,HM2,HM3}. In particular we mention a specific model of HM in which the cutoff length is proportional to the Ricci scale\cite{R1,V}, see also Ref.\cite{V1}. In relation to the study of the dynamics of perturbations, this was analyzed in Refs.\cite{V0,V2}. The appearance of instabilities in the dark sector, through the perturbative dynamics, occurs when the EoS parameter $w$ crosses the value $w=-1$ in models with a dynamical EoS. In this context, the study of this crossing of the EoS parameter and the appearance of instabilities in the perturbative dynamics for non-interacting Ricci holographic model, was performed in Ref.\cite{DelCampo2013}. In the present paper we study the background dynamics and also present the analysis of linear perturbations in the framework of gauge-invariant variables in comoving gauge, for the interacting dark sector, identifying the source of instabilities. In particular we consider that the dark energy density corresponds to the holographic Ricci DE model, and then we extended the study developed in Ref.\cite{DelCampo2013}, but now considering an interacting dark sector. Here we study the background equations, the linear perturbations and the appearance of instabilities. In order to evade these instabilities we develop an appropriate holographic Ricci interacting dark energy model and we find drastic changes on the constraints of the parameters. This article is organized as follows. In Sect. II we present the background equations for the interacting dark sector. In Sect. III we analyze the dynamics of perturbations in a general framework for a single fluid and two interacting fluids. In Sect. IV we study the evolution of linear perturbations and identify the sources of instabilities. In Sect. V we consider a specific holographic dark energy model, known as Ricci DE. Here we study the background dynamics and analyze the observational tests on this model, considering the SNIa and $H(z)$ data sets. In Sect.VI we study the linear perturbations and identify the instabilities in our interacting model. We also analyze the high-redshift limit of our model and we compare it with $\Lambda$CDM model. In Sect. VII we study how to avoid these instabilities and develop an appropriate model . Finally, Sect. VIII summarizes our results and exhibits our conclusions.
\begin{figure}[htb] \includegraphics[]{Fig7.eps} \caption{The evolution of the perturbation $\delta^{cm}_m$ as function of the scale factor $a$, for the scale $k=1.5 h^{-1} Mpc^{-1}$. Here we have used the values $\beta_1=-0.05$, $\beta_2=0.18$ and $w_0=-0.95$ for the model 1, and $\beta_1=-0.03$, $\beta_2=0.22$ and $\Omega_{m0}=0.25$ for the model 2.}\label{fig8} \end{figure} In this paper we have analyzed an interacting model of dark energy and dark matter in order to describe of late cosmic acceleration of the universe. Under a general formalism we have described the perturbative dynamics for these two interacting fluids. In this general analysis we have considered the timelike part of the balance equation, the momentum balance and the momentum transfer $Q^{\alpha}$ associated to the interaction term $Q$. From the functions of contrast for dark matter and dark energy we have studied the perturbative dynamics considering a gauge-invariant treatment in comoving gauge. On the other hand, we have obtained the total non-adiabatic and dark energy pressure perturbations and we found the relation between these gauge-invariant quantities. Also, we have identified the sources of instabilities in DE models with a dynamical EoS parameter that presents a phantom crossing. As a concrete example we have considered an interaction term $Q$ between the holographic Ricci-DE and DM. Here we have studied that the interaction term $Q$, depends on the energy densities of both components multiplied by a quantity with units of the inverse of time (proportional to the Hubble parameter). From the background equations we have obtained the constraints on the parameters characterizing the interaction by considering the observational analysis from the SNIa and $H(z)$ tests. Here from the background dynamics we have found that the best-fit values for the parameters of the interaction are $\beta_1=-0.05^{+0.05+0.08+0.10}_{-0.05-0.07-0.09}$, $\beta_2=0.18^{+0.04+0.06+0.08}_{-0.04-0.06-0.08}$, and for the EOS parameter $w_0=-0.95^{+0.05+0.06+0.08}_{-0.05-0.07-0.09}$. In our perturbative analysis we have found that, in this best-fit model, the instabilities appear at the moment when the EoS parameter $w$ crosses the value $w\sim -1$ and we have noted that this feature does not depend on the interaction term $Q$. In order to avoid these instabilities in the perturbative analysis and develop an appropriate model for any finite time, we have obtained a specific value of the scale factor denoted as $a_i$. From this value of $a_i$ we have obtained two independent conditions to avoid the instabilities in our specific model, namely: i) $Cr_0=-D$ and ii) $r_0=3(w_0+1)$. Considering the first condition we have found that if $w_0<1+r_0$, the model is disproved from observations, since under this requirement the model does not present an accelerate scenario. Otherwise if $w_0>1+r_0$, we have obtained an accelerate phase, however this condition corresponds to a phantom model, nothing that the instabilities take place in the past time. In this way, we have obtained that the first condition is not suitable. From the second condition i.e., $r_0=3(w_0+1)$, we have found an accelerate phase of the universe and also corresponds to an appropriate model for any finite time. In order to avoid the instabilities in the perturbative dynamics, we have noted that this result agrees with obtained in Ref.\cite{DelCampo2013} and becomes independently of the interaction term. Moreover we have obtained that the constraint on the Ricci parameter $c^2=1/2$ is fixed for the second condition, since as from background $c^2=2(r_0-3w_0+1)^{-1}$ and together the second condition $r_0=3(w_0+1)$, then $c^2=1/2$, independency of the values $r_0$ , $w_0$ and the energy transfer rate $Q$. Also, from this condition we have found, in order to have an appropriate model, a new sets of best-fit values for the interaction parameters given by $\beta_1=-0.03^{+0.04+0.07+0.08}_{-0.10-0.14-0.17}$, $\beta_2=0.22^{+0.09+0.14+0.18}_{-0.10-0.14-0.17}$, and $\Omega_{m0}=0.25^{+0.01+0.03+0.04}_{-0.01-0.02-0.03}$. Here we have observed a drastic change in the values of the parameters $\beta_1$ and $\beta_2$ in order to avoid the singularity from the perturbative dynamics. For the parameter $\beta_1$ we have found that the increased is the order of 40$\%$ and for $\beta_2$ is the order of 22$\%$. Finally, we would like to point out that in models that have a phantom crossing, and in particular for the holographic models (with and without interaction), it is necessary to be cautious when only the background level observational tests are being considered. Here, we have shown that in spite of a good agreement with data and an adequate background dynamic, this could lead to inviable models at perturbative level.
16
7
1607.01806
In this paper we investigate the consequences of phantom crossing considering the perturbative dynamics in models with interaction in their dark sector. By mean of a general study of gauge-invariant variables in comoving gauge, we relate the sources of instabilities in the structure formation process with the phantom crossing. In order to illustrate these relations and its consequences in more detail, we consider a specific case of an holographic dark energy interacting with dark matter. We find that in spite of the model is in excellent agreement with observational data at background level, however it is plagued of instabilities in its perturbative dynamics. We reconstruct the model in order to avoid these undesirable instabilities, and we show that this implies a modification of the concordance model at background. Also we find drastic changes on the parameters space in our model when instabilities are avoided.
false
[ "dark matter", "models", "instabilities", "phantom crossing", "background level", "background", "an holographic dark energy", "interaction", "their dark sector", "comoving gauge", "observational data", "its perturbative dynamics", "the perturbative dynamics", "the concordance model", "excellent agreement", "these undesirable instabilities", "more detail", "the phantom crossing", "the structure formation process", "our model" ]
11.258311
0.570833
89
12464956
[ "Bianchini, Federico", "Renzi, Alessandro", "Marinucci, Domenico" ]
2016JCAP...11..050B
[ "Needlet estimation of cross-correlation between CMB lensing maps and LSS" ]
2
[ "Astrophysics Sector, SISSA, Via Bonomea 265, I-34136 Trieste, Italia ; INAF, Osservatorio Astronomico di Trieste, via Tiepolo 11, 34131, Trieste, Italy; INFN, Sezione di Trieste, Via Valerio 2, I-34127 Trieste, Italy;", "Dipartimento di Matematica, Università di Roma Tor Vergata, Via della Ricerca Scientifica 1, 00133 Roma, Italia ; INFN, Sezione di Roma 2, Università di Roma Tor Vergata, Via della Ricerca Scientifica 1, 00133 Roma, Italia;", "Dipartimento di Matematica, Università di Roma Tor Vergata, Via della Ricerca Scientifica 1, 00133 Roma, Italia ; INFN, Sezione di Roma 2, Università di Roma Tor Vergata, Via della Ricerca Scientifica 1, 00133 Roma, Italia;" ]
[ "2019MNRAS.482.2670B", "2019MNRAS.485.1339B" ]
[ "astronomy" ]
8
[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1953ApJ...117..134L", "1997PhRvD..55.5895T", "2000ApJ...538..473L", "2002ApJ...566...19K", "2002ApJ...567....2H", "2004MNRAS.349..603E", "2005ApJ...622..759G", "2006MNRAS.365..891V", "2006PhR...429....1L", "2006PhRvD..74d3524P", "2007MNRAS.376.1211M", "2007PhRvD..76d3510S", "2008EJSta...2..332L", "2008MNRAS.383..539M", "2008PhRvD..77l3520G", "2008PhRvD..78l3533G", "2009AnSta..37.1150B", "2009ApJ...701..369R", "2010MNRAS.402L..34P", "2010PASP..122..499E", "2010arXiv1006.5344D", "2011arXiv1110.3193L", "2012ApJ...755...19D", "2012ApJ...761..152T", "2012MNRAS.427.3044S", "2013SPIE.8858E..0IM", "2014A&A...571A..12P", "2014A&A...571A..19P", "2014IAUS..306...48T", "2014MNRAS.442..821M", "2014PhRvD..90j3532D", "2015ApJ...802...64B", "2015JCAP...01..013R", "2015JCAP...03..049C", "2015MNRAS.451..849A", "2015PhRvD..92l3010L", "2015arXiv150203120L", "2016A&A...594A...1P", "2016A&A...594A...9P", "2016A&A...594A..17P", "2016A&A...594A..21P", "2016ApJ...825...24B", "2016ApJS..225....5B", "2016MNRAS.456.3213G", "2016MNRAS.460..434H", "2016MNRAS.460.3014R", "2016MNRAS.461.4099B", "2016MNRAS.463.2310R" ]
[ "10.1088/1475-7516/2016/11/050", "10.48550/arXiv.1607.05223" ]
1607
1607.05223_arXiv.txt
\label{sec:intro} One of the main puzzles of modern cosmology is the understanding of the mechanism that sources the late-time accelerated expansion of the Universe. Whether it is associated to an exotic form of energy or to some modifications of general relativity, the different scenarios can only be disentangled by probing the perturbations evolution over cosmic time. In this context, galaxy clustering and weak gravitational lensing have become promising probes not only to investigate cosmic acceleration but also the dark matter and neutrino sectors. While the analysis of the data from the Planck satellite is approaching to an end, yielding a breakthrough in many respects for what concerns CMB studies \cite{PlanckCollaboration2015c}, such fundamental issues have triggered the upcoming experimental efforts and in the next few years galaxy surveys such as the European Space Agency's (ESA) satellite Euclid\footnote{\url{http://sci.esa.int/euclid/}} \cite{Laureijs2011}, the Dark Energy Spectroscopic Instrument (DESI\footnote{\url{http://desi.lbl.gov}}), the Large Synoptic Survey Telescope (LSST\footnote{\url{http://www.lsst.org}}) and the Wide Field Infrared Survey Telescope (WFIRST\footnote{\url{http://wfirst.gsfc.nasa.gov}}), along with a plethora of ground-based high-sensitivity CMB experiments like the Simons Array\footnote{\url{http://cosmology.ucsd.edu/simonsarray.html}}, the South Pole Telescope (SPT-3G)\footnote{\url{https://pole.uchicago.edu/spt/}}, and the Advanced Atacama Cosmology Telescope (AdvACT)\footnote{\url{https://act.princeton.edu}}, will carry out observations devoted to shed light on the physics behind the dark components. In these experiments, operating and under design and construction towards the efforts of the next decade (including ground-based facilities such as the Simons Observatory\footnote{\url{https://simonsobservatory.org}} and CMB-S4, as well as the proposed space satellites COrE\footnote{\url{http://www.core-mission.org}} e LiteBIRD\footnote{\url{http://litebird.jp/eng/}}), the role of CMB-LSS cross correlation is double: on one side, yielding constraints on dark energy and matter through the analysis of CMB lensing by forming LSS, and on the other, to de-lens the B-modes of polarization in order to improve the constraint, or measure, of the power from primordial gravitational waves. In particular, LSS data gathered from Euclid in the form of weak lensing and galaxy catalogues will provide an excellent tracer for the underlying gravitational potential which is responsible for the CMB lensing effect. It is then only natural to cross-correlate CMB lensing maps with LSS data to improve the constraints on dark energy models and cosmological parameters, similarly to what has been done with CMB temperature and LSS maps in order to extract faint large scale signal like the integrated Sachs-Wolfe effect (iSW), see for instance \cite{Pietrobon2006, Vielva2006, McEwen2007, Munshi:2014tua,Ade:2013dsi,Ade:2015dva}. CMB lensing-galaxy cross-correlation measurements have found different applications in cosmology, such as the reconstruction of the galaxy bias redshift evolution \cite{Bianchini2015,Allison2015a,Bianchini2016}, the investigation of the growth of structures \cite{Giannantonio2016}, and the augmentation of the absolute cosmic shear calibration \cite{Baxter2016}. All analyses reported to date have reconstructed the 2-point statistics either in harmonic or real space. The optimal power spectrum estimator in harmonic space for auto and cross-correlation in presence of mask and anisotropic noise is well known \cite{Tegmark1997} and was used for cross-correlation analysis in \cite{Smith2007} and \cite{Schiavon2013}. Modern iterative algorithms make the exact optimal estimation computationally feasible; despite being potentially suboptimal in cases of very small $f_{\rm sky}$ and highly non-uniform noise, the computational convenience of a fast pseudo-$C_\ell$ (PCL) estimator remains an important property, especially when a cross-correlation analysis must be implemented on a variety of different masks, due to different observational strategies in multiple experiments. This is especially relevant when cross-correlating lensed CMB maps with LSS data. Note that most analysis with very small $f_{\rm sky}$ have been so far performed in the flat-sky approximation, in conjunction with the MASTER algorithm, and they provide a nearly optimal power spectrum estimation. In this paper, we shall use instead a procedure based on a wavelet-domain approach; more precisely, we shall discuss how to modify the PCL algorithm to perform a needlet cross-correlation analysis. Since needlet transform is linear in the data, it cannot perform better than the optimal estimator; however it can improve the performance of linear estimators in the presence of masks, as we shall discuss below. At the same time, a needlet estimator mantains the computational convenience of a nearly-optimal PCL estimator, provided that the noise properties are fairly uniform. As discussed in many previous references, needlets are a form of spherical wavelets which were introduced in functional analysis and statistics by \cite{Narcowich2006,Baldi2009a} and have then found a number of different applications in the cosmological community over the last decade; we recall for instance \cite{Marinucci2007} for a general description of the methods, \cite{Lan2008,Rudjord2009a,Pietrobon2010a,Donzelli2012,Regan2015,Ade:2015ava} for non-Gaussianity estimation, \cite{Delabrouille2010,2014A&A...571A..12P,Adam:2015tpy,Rogers2016,Rogers2016a} for foreground component separation, \cite{Geller2008,Leistedt2015a,Ade:2015ava} for polarization data analysis, \cite{Durastanti2014,Leistedt2015} for extension in 3d framework and \cite{Troja2014,Regan2015} for trispectrum analysis. The advantages of needlets, like those of other wavelets system, have been widely discussed in the literature; in short, they are mainly concerned with the possibility to exploit double localization properties, in the real and harmonic domain. Despite this localization in the real domain, we show here that the performance of a needlet cross-correlation estimator deteriorates badly in the presence of very aggressive sky-cuts (i.e., experiments with sky coverage much smaller than 50\%). In this paper, we show how the performance of this estimator can be restored by a MASTER-like correction. Thus achieving signal-to-noise figure of merits which are in some aspect superior to the corresponding results for power spectrum methods; the terms of this comparison are explained in more details below. The plan of the paper is as follows. In section \ref{sec:theo} we review quickly some background material on both harmonic and needlet cross-correlation analysis; we then proceed in section \ref{sec:master} to introduce the MASTER-like algorithm for the needlets cross-correlation estimator. Numerical evidence and some comparison on the performance of these procedures are collected in section \ref{sec:num_ev}, while final considerations are presented in section \ref{sec:conclusion}.
\label{sec:conclusion} Cross-correlation analyses between independent cosmological datasets have the advantage to be potentially immune to any known (and unknown) systematics, as well as to extract signals hidden in noisy data. In this way, cross-correlation measurements can provide us with a clearer view of the large scale distribution of matter, fundamental to reconstruct the dynamics and the spatial distribution of the gravitational potential that can be then translated into constraints on cosmological parameters, breaking degeneracies with the astrophysical ones. In this paper we begin a systematic analysis of the scientific potential associated to the expansion of the analysis domain in CMB-LSS cross-correlation studies to include the localization in the harmonic and spatial domains. In this initial application, by exploiting an ensemble of simulations, we have shown that under the same observational configurations the needlet spectral estimator can enjoy some advantages over the harmonic one, thanks to the excellent needlets localization properties in both pixel and frequency space, as well as their optimal window function. Moreover, we have completed an initial needlet based analysis pipeline throughout the implementation of a novel MASTER-like approach for needlet spectral reconstruction in the case of aggressive masking ($f_{\rm sky}\simeq 0.01$), reporting an higher total $S/N$ with respect to its harmonic counterpart. As we discussed earlier, these comparisons must be considered with some care, because the bin size is intrinsically different in the harmonic and needlet cases. Motivated by these positive indications and results, in future research we plan to explore further the trade-off between $S/N$ and multipole localization, so as to achieve optimal bandwidth selection for a given experimental setting (such as the Euclid coverage mask). We also aim at applying this machinery to accurate CMB maps lensed with ray-tracing techniques \cite{Calabrese2015} and realistic galaxy mock catalogues based on N-body simulations by adopting, on the CMB side, the projected accuracy and sensitivity of forthcoming polarization oriented CMB probes, targeting the B-modes from cosmological gravitational waves and gravitational lensing. This work is of course preparatory for application to real data, from currently available LSS maps such as Herschel and WISExSCOS Photometric Redshift Catalogue (WISExSCOSPZ) \cite{Bilicki2016} to upcoming surveys such as Euclid, LSST, DESI, and WFIRST, in order to robustly extract cosmological information from cross-correlation measurements.
16
7
1607.05223
In this paper we develop a novel needlet-based estimator to investigate the cross-correlation between cosmic microwave background (CMB) lensing maps and large-scale structure (LSS) data. We compare this estimator with its harmonic counterpart and, in particular, we analyze the bias effects of different forms of masking. In order to address this bias, we also implement a MASTER-like technique in the needlet case. The resulting estimator turns out to have an extremely good signal-to-noise performance. Our analysis aims at expanding and optimizing the operating domains in CMB-LSS cross-correlation studies, similarly to CMB needlet data analysis. It is motivated especially by next generation experiments (such as Euclid) which will allow us to derive much tighter constraints on cosmological and astrophysical parameters through cross-correlation measurements between CMB and LSS.
false
[ "CMB needlet data analysis", "cross-correlation measurements", "CMB", "CMB-LSS cross-correlation studies", "LSS", "CMB-LSS", "cosmic microwave background", "masking", "different forms", "maps", "(LSS) data", "large-scale structure", "the operating domains", "Our analysis", "next generation experiments", "correlation", "noise", "cosmological and astrophysical parameters", "-", "Euclid" ]
12.746417
1.839821
103
12515390
[ "Pawar, P. K.", "Dewangan, G. C.", "Patil, M. K.", "Misra, R.", "Jogadand, S. K." ]
2016RAA....16..169P
[ "On the reality of broad iron L lines from the narrow line Seyfert 1 galaxies 1H 0707-495 and IRAS 13224-3809." ]
1
[ "-", "-", "-", "-", "-" ]
[ "2022MNRAS.516.2374P" ]
[ "astronomy" ]
4
[ "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.48550/arXiv.1607.02635" ]
1607
1607.02635_arXiv.txt
Active galactic nuclei (AGN) are thought to be powered by the accretion of matter onto the central super massive black hole (SMBH). The surrounding matter forms an optically thick, geometrically thin disk that radiate mainly in the optical/UV region. The broadband \xray{} spectrum of AGN follows a power law shape and is thought to originate from the Compton up-scattering of low energy disk photons by the relativistic electron cloud present in the hot Comptonizing corona \citep{1976ApJ...204..187S,1985ApJ...289..514Z,1980A&A....86..121S,1991ApJ...380L..51H}. The geometry and origin of this corona is, however, still unclear. In addition to this primary power law continuum several other features are also apparent, which include reflection hump in the energy range 10--50 \kev{}, broad and skewed Fe \kalpha{} fluorescent line around 6.4 \kev{} and soft excess emission below 1 \kev{}. Origin of the reflection hump and the fluorescent Fe \kalpha{} line is generally attributed to the reflection of the power law photons by the relatively cold accretion disk \citep{1991MNRAS.249..352G}. First clear evidence for the presence of the extremely broad, skewed iron line came from the {\it ASCA} long observation of MCG--6--30--15 \citep{1995Natur.375..659T}. Similar line profiles were later noticed in several other AGN (for a review, see \citealt{2007ARA&A..45..441M} and references therein). Detection of the broad Fe \kalpha{} is important as it carries signature of the inner accretion disk close to the SMBH e.g. the line energy tells us the ionization state of the disk, the inner radius of accretion disk can be inferred from the redward wing of the line which extends sometimes even down to $\sim$ 2 \kev{}. The line also provides an unique tool to test and verify the GR theory \citep{2000PASP..112.1145F}. As mentioned earlier, the AGN spectrum consists of features like soft excess emission, reflection hump and broad Fe K lines. These features can be reproduced using the blurred reflection of primary continuum from partially ionized accretion disk \citep{1991MNRAS.249..352G}. Several versions of self--consistent reflection models are available e.g. $reflionx$, $relxill$ etc. which provide good approximation to the observed data \citep{2005MNRAS.358..211R,2006MNRAS.365.1067C,2014ApJ...782...76G}. Assuming lamp--post geometry, the recent blurred reflection models describe both the spectral shape and the observed variability of iron lines \citep{2003MNRAS.344L..22M,2004MNRAS.353.1071F,2010MNRAS.401.2419Z,2014MNRAS.443.1723P,2015MNRAS.446..633G}. The AGN \xray{} spectra, however, can also be fitted using complex partial--covering absorption model which require the central engine to be obscured by complex absorbing clouds/zones having varying column density, covering fraction and ionization. Here, the observed spectral variability is attributed to the motions of these absorbing clouds \citep{2008A&A...483..437M,2009A&ARv..17...47T,2014PASJ...66..122M,2014ApJ...787...83M,2014MNRAS.441.1817P}. Both the blurred reflection and partial covering absorption models provide statistically comparable spectral fits and hence must be judged based on the best fit spectral parameter values. One of the key fitting parameter in reflection scenario is the iron abundance. Observationally, many AGN require over abundance of iron relative to the solar value e.g. Ark 120 \citep{2014MNRAS.439.3016M}; NGC 1365, \citep{2010MNRAS.408..601W}; \hhh{}, \citep{2009Natur.459..540F}. This over abundance of iron must also produce the accompanying Fe \lalpha{} which has not been observed in those AGN spectra. These anomalies cast shadows on the reality of the iron lines and hence on the reflection scenario. Recently, the accompanied Fe \lalpha{} line was detected in two extreme cases of typical narrow line Seyfert 1 galaxies (NLS1) \hhh{} and \iras{} using \xmm{} data by \cite{2009Natur.459..540F} and \cite{2010MNRAS.406.2591P}, respectively. In fact, \cite{2009Natur.459..540F}, using Fe \lalpha{} line claimed a lag of 30$s$ between the direct power law component (1--4 \kev{}) and the reflection component (0.3--1 \kev{}), as it provided better statistics compared to the \kalpha{} line. However, the AGN spectrum can be complex due to the presence of multi--component absorber and soft \xray{} excess and strong spectral variability may introduce artificial spectral features in the mean spectrum. Artifacts in the AGN spectra introduced due to the spectral modelling have been reported earlier e.g. the soft excess is an artifact of the deficit of emission due to smeared wind absorption while the complex partial-covering absorption model can mimic the broad Fe \kalpha{} line quite well. Both these artifacts are model dependent. The \xray{} spectrum of NLS1 is variable and can vary within few $ks$ \citep{1996A&A...305...53B}. In the case of \xmm{} observations, where typical exposure time is $>$ 100 $ks$, it is very likely that both the flux and the spectral shape may change significantly. These changes in the flux and/or spectral parameters may introduce the broad bumps/dips in the time averaged spectrum. In this paper, using a two component phenomenological model and using time resolved spectroscopy, we try to comment on the reality of Fe \lalpha{} line observed in these two AGN. The observational details and data reduction techniques are summarized in section 2. The spectral analysis and the results are presented in section 3. The discussion and conclusion is given in section 4.
We performed time resolved spectroscopy of NLS1 galaxies \hhh{} and \iras{} using $\sim$ 100 $ks$ long \xmm{} observations. NLS1 galaxies are generally characterized by steep power law spectrum and rapid variability. We found that both these AGN show strong short--term variations within the observations. Figure~\ref{lc_variations} shows the significant variability observed in the soft band (0.3--1 \kev{}), the hard band (1.5--5 \kev{}) and in the hardness ratio. The variability in hardness ratio suggests the spectral variability within the observation. These spectral variations motivated us to study the variability of spectral components and the artifacts which could be introduced by such variability in the average spectrum. In this study we performed time resolved spectroscopy by generating spectra from multiple small segments. In Fig.~\ref{spec_examples} we show unfolded spectra of two typical segments fitted with a $power$ $law$ model (with $\Gamma$ fixed at 2) to show the short--term spectral variations in both the sources. This figure reveals the significant spectral variability within the observation. We find that each segment spectra can be easily modelled using two component model (PL+BB) modified by the Galactic absorption. As expected, our spectral result exhibit that all the spectral parameters were variable and are plotted in Fig.~\ref{spec_variations}. Variability of $\Gamma_{PL}$ and kT$_{BB}$ indicates that the observed variability is not just due to the flux variations but the spectral shape is also changing.\\ The best fit model of each spectra were used to simulate the data which were co--added to get the combined simulated spectra. These combined simulated spectra were later compared with the actual time averaged spectra. The result of this spectral fitting is tabulated in Table~\ref{spec_res}. The spectral parameters show variations between the simulated and average spectra, which is solely because the GTI of simulated spectrum is a subset of total exposure of the time averaged spectrum. The extra time in the average spectra caused the variation in the spectral parameters seen in Table~\ref{spec_res}. We find that the time averaged spectra showed significant deviations and required additional line components, therefore we added two $Laor$ lines to the model. In simulated spectra we do not see line at 6.4 \kev{} as our initial model did not have the line model, however, no deviation was seen near 0.9 \kev{}. This suggest that the line we see in the time averaged spectra is not the artifact of the variation of spectral components. In fact, no deviation near 0.9 \kev{}, this is an independent way of proving that the line indeed is a genuine feature. Even if we did not find line feature but a positive deviation around 0.9 \kev{} would have certainly lowered the overabundance required in these objects and would have improved the current reflection models.
16
7
1607.02635
We performed time resolved spectroscopy of 1H0707-495 and IRAS 13224-3809 using long XMM-Newton observations. These are strongly variable narrow line Seyfert 1 galaxies and show broad features around 1 keV that has been interpreted as relativistically broad Fe L$\alpha$ lines. Such features are not clearly observed in other AGN despite sometimes having high iron abundance required by the best fitted blurred reflection models. Given the importance of these lines, we explore the possibility if rapid variability of spectral parameters may introduce broad bumps/dips artificially in the time averaged spectrum, which may then be mistaken as broadened lines. We tested this hypothesis by performing time resolved spectroscopy using long (&gt; 100 ks) XMM-Newton observations and by dividing it into segments with typical exposure of few ks. We extracted spectra from each such segment and modelled using a two component phenomenological model consisting of a power law to represent hard component and a black body to represent the soft emission. As expected both the sources showed variations in the spectral parameters. Using these variation trends, we simulated model spectra for each segment and then co-added to get a combined simulated spectrum. In the simulated spectra, we found no broad features below 1 keV and in particular no deviation near 0.9 keV as seen in the real average spectra. This implies that the broad Fe L? line that is seen in the spectra of these sources is not an artifact of the variation of spectral components and hence providing evidence that the line is indeed genuine.
false
[ "broadened lines", "broad features", "spectral components", "few ks", "spectral parameters", "model spectra", "hard component", "relativistically broad Fe L$\\alpha$ lines", "segments", "typical exposure", "averaged spectrum", "rapid variability", "Such features", "long XMM-Newton observations", "high iron abundance", "the best fitted blurred reflection models", "variations", "the broad Fe L? line", "broad bumps/dips", "Fe L" ]
15.903364
8.059074
118
12586165
[ "Aldi, Giulio Francesco", "Bozza, Valerio" ]
2017JCAP...02..033A
[ "Relativistic iron lines in accretion disks: the contribution of higher order images in the strong deflection limit" ]
19
[ "Dipartimento di Fisica \"E.R. Caianiello\", Via Giovanni Paolo Secondo 132, Fisciano, SA, I-84084 Italy ; Istituto Nazionale di Fisica Nucleare, Sezione di Napoli, Via Cintia, Napoli, 80126 Italy;", "Dipartimento di Fisica \"E.R. Caianiello\", Via Giovanni Paolo Secondo 132, Fisciano, SA, I-84084 Italy ; Istituto Nazionale di Fisica Nucleare, Sezione di Napoli, Via Cintia, Napoli, 80126 Italy;" ]
[ "2017FTP...189...63H", "2017IJMPD..2641013B", "2017JCAP...07..045C", "2021JCAP...10..013I", "2021PhRvD.104f4022T", "2021PhRvD.104l4016T", "2021arXiv210707146T", "2021arXiv210900495T", "2022EPJC...82.1175F", "2022JCAP...11..006G", "2022PhRvD.105b4009T", "2022PhRvD.105f4013T", "2022PhRvD.105h4036T", "2022PhRvD.106f4033T", "2022PhRvD.106h4025T", "2022arXiv220810197T", "2023EPJC...83..284T", "2023FrP....1113909C", "2023PhRvD.108l4054D" ]
[ "astronomy" ]
3
[ "Astrophysics - High Energy Astrophysical Phenomena", "General Relativity and Quantum Cosmology" ]
[ "1959RSPSA.249..180D", "1968PhRv..174.1559C", "1985MNRAS.216P..65B", "1989MNRAS.238..729F", "1991ApJ...376...90L", "1991MNRAS.250..629K", "1992ApJ...400..163B", "1993ApJ...413..680H", "1994ApJ...425...63B", "1994ApJ...435...55B", "1995Natur.375..659T", "1996ASPC..101...17A", "2000MNRAS.312..817M", "2002A&A...383L..23M", "2002PhRvD..66j3001B", "2003PhRvD..67j3006B", "2004ApJS..153..205D", "2004GReGr..36..435B", "2004MNRAS.352..353B", "2005MNRAS.363..177C", "2005PhRvD..72h3003B", "2005astro.ph..7409F", "2006PhRvD..74f3001B", "2007CQGra..24S.259N", "2007GReGr..39.1563I", "2007PhRvD..76h3008B", "2008MNRAS.386..759N", "2008PhRvD..78f3014B", "2009A&A...507....1S", "2009ApJ...696.1616D", "2009ApJ...703L.142D", "2010ApJ...718..446J", "2010GReGr..42.2269B", "2010PhDT.......467S", "2012ApJ...758...30J", "2013arXiv1306.2334M", "2016ApJ...826...87W" ]
[ "10.1088/1475-7516/2017/02/033", "10.48550/arXiv.1607.05365" ]
1607
1607.05365_arXiv.txt
\label{sec:intro} Black holes reveal their presence in the universe thanks to diffuse material in orbital motion around them that emits electromagnetic radiation particularly in the X-ray band. The emission is attributed to the radiation of inward-spiralling matter in the form of an accretion disk. The dominant features in the X-ray spectrum seen by a distant observer are the iron lines with broad skewed profiles. The energy of the Fe K$\alpha$ line is about 6.4 KeV and the natural bandwidth is a few eV. Relativistic effects and the enormous velocity of the material in the disk are the principal causes of the line deformation \cite{fabian1989} and the typical observed bandwidth are about thousand of KeV \cite{Tanaka1995,Martocchia2002}. Furthermore, it was noted that Fe K$\alpha$ is unaffected by recombination phenomena across the geodesics path near the accretion disk environment. This underlines the reality and the importance of the broad red wing of the line \cite{fabian2008} as a tool to analyze the main characteristics of the central black hole. Fabian et al. constructed a \emph{diskline} model for the relativistic line around a non-rotating black hole \cite{fabian1989}. A model of the relativistic line for a maximally rotating Kerr black hole, based on numerical approach, was developed independently by Kojima \cite{kojima1991}, Laor \cite{laor91} and Arnaud \cite{arnaud1996}. This allowed to extract unique information on the structure, geometry and dynamics of the accretion flow in the immediate vicinity of the central black hole \cite{fabian1989,fabian2008,laor91}. A. Martocchia, V. Karas and G. Matt developed a detailed model of the relativistic effects on both the reflection continuum and the iron line profile \cite{martocchia2000}. M. Dovciak, V.Karas and T. Yaqoob analyzed X-ray spectra of black hole accretion disk by introducing new routines for the XSPEC package \cite{arnaud1996,dovciak2004}. Similarly, Cadez and Calvani developed an advanced code including warped disks and other effects \cite{Cadez2005}. We also mention the fast code developed by Dexter and Agol \cite{DexAgo,DexAgoFra}. In these codes, the photons coming from the relativistic accretion disk are followed along their geodesics in Kerr space using semi-analytical integration in terms of elliptic integrals. As well known, gravitational lensing is responsible for the creation of multiple images of the same source. In the case of an accretion disk, we will have a direct image from photons coming from the disk side face to the observer, but also a first order image from photons emitted from the opposite side and turning around the black hole. As shown by G. Bao, P. Hadrava and E. \O stgaard, the first order image can contribute to the total flux as much or even more than the direct image \cite{bao1994}. Also it has been shown in Ref. \cite{fuerst} that the most significant contribution to the total flux appears for higher inclination of the accretion disk with respect to the line of sight. Higher order images, generated by photons performing one or more complete revolutions around the black hole, are exponentially suppressed. However, their total contribution can reach a few percents of the total flux and become dominant at their peak frequencies \cite{bao1994}. The problem with higher order images is that they pose the most sever computational challenge for numerical codes. So, apart from the Schwarzschild case, illustrated in detail in Refs. \cite{bao1994,bao1994a,bao1992} using elliptic integrals, higher order images are often neglected. Recently, Johannsen and Psaltis have shown that the higher order images of the accretion disk would build up a bright ring surrounding the shadow, which can be used to estimate the ratio $M/D$ of the black hole \cite{JohPsa}. Perspectives for detecting this feature by VLBI are currently discussed \cite{Joh2012}. Once we are convinced of the importance of describing higher-order images of disk accretions correctly, in this paper, we propose to apply the well-known approximation technique of the Strong Deflection Limit (SDL), which has been successfully employed in a vast literature to describe the formation of higher order images of isolated sources in a large number of metrics (see e.g. \cite{Bozza2010} and references therein). The Schwarzschild case was considered by Darwin \cite{Darwin1959} as a limit of the exact formula of the deflection angle based on elliptic integrals. The method was then generalized in Ref. \cite{Bozza2002} to any spherically symmetric metrics, and then to rotating black holes and to sources in arbitrary positions \cite{Bozza2003,Bozza2005,Bozza2006,Bozza2007}. Higher orders in the expansion were calculated in Ref. \cite{Iyer2007}. The time delay between different higher order images was derived in Ref. \cite{BozzaMancini2004}. It is interesting to note that alternatives to the Schwarzschild solution arising in several gravitational theories can be typically distinguished by their higher order images pattern, which strictly depends on the derivatives of the metric coefficients, as shown in Ref. \cite{Bozza2002}. In the framework of the study of the iron emission lines created in accretion disks, these analytical approximations can be very useful to set up a simple and accurate treatment of the contribution of the higher order images to the spectral lines. Coupled with a numerical treatment for the direct and first order images, this treatment can provide the missing complement to build up the whole picture of the relativistic effects affecting the formation of such high energy emission lines. In this paper we will achieve this goal through the following steps. In Section 2 we specify the model for a thin accretion disk used thereafter. Section 3 introduces the strong deflection limit formalism on which our method is based. We then calculate all needed functions fully analytically. Section 4 contains a thorough discussion of the physical appearance of higher order spectral contributions for various ranges of the parameters, supplying several interesting analytical formulae describing peaks, ranges and other features. We discuss the relation to other works and future perspectives in Section 5.
In this paper, we have proposed to use the strong deflection limit formalism to obtain accurate analytical results on the higher order images of the accretion disk. We have derived the magnification, redshift and emission angle functions in a fully analytical way. These functions enter the numerical integration where the radial and angular emissivity profiles can be specified according to our preferred disk model. We have shown that two different regimes exist for the line shape. For edge-on observers ($\mu_o<1/\sqrt{2}$, corresponding to $\vartheta_o<45^\circ$), the peak frequency is determined by a saddle point in the redshift function. For face-on observers ($\mu_o>1/\sqrt{2}$), the peak frequency is determined by the redshift at the ISCO. All these results do not depend on the model of accretion disk, because all the quantities only depend on the metric of space-time. Indeed, the ansatz on the intrinsic emissivity enters only in the numerical integration to determine the profile shape. So, our formalism can be applied to any physically motivated models of disks. In the literature there are many works about the relativistic shapes of the Fe K$\alpha$ line. Some of them also try to address the role of higher order images of the accretion disk, but the results seem quite discrepant. Ref. \cite{beck} uses a full numerical approach to determine the shape of second-order line profile. Their profiles for the higher order contribution are dramatically different from those we have elaborated, with the presence of several peaks and more complex structures. On the other hand, we find a satisfying agreement with the results obtained by Bao et al. \cite{bao1994} where we can see that the structure of spectra for a quasi-equatorial observer is very similar to the one determined by us. We have been able to identify the origin of the lonely peak appearing in the contribution by higher order images, attributing it to the existence of a saddle point in the redshift function. The position of the peak is tracked by the redshift of this saddle point. The fact that we are proposing an approximation in the deflection angle (the strong deflection limit) so as to obtain simple analytical formulae has no effect on the overall structure, which is determined by the redshift function. The approximation intervenes in the calculation of the magnification, where it is accurate at the percent level, for the second and higher order images. Therefore, the structure of the contribution by higher order images appears relatively simple and we cannot imagine any physical origin for the multi-peak structures appearing in Ref. \cite{beck}, which might be the fruit of numerical errors. The existence of other independent calculations agreeing with our results \cite{bao1994} supports this conclusion. These examples show how higher order images can be tricky to model in numerical codes for relativistic emission lines. Although they contribute by less than $1\%$ to the total flux in the line, they should not be overlooked at all, since at specific frequencies (namely, their peak frequency) they might give rise to localized bumps in the total line profile, as we can see in Fig. \ref{fig:ratio2}. For relatively bright sources such as Cygnus X-1, an accuracy of $1\%$ has already been reached \cite{Walton2016}. So, models trying to fit the observations should not neglect the contribution of higher order images. In this work we have confined our investigation to the Schwarzschild black hole, in order to obtain simple and appealing analytical formulae. However, the strong deflection limit formalism has been extended to slowly rotating black holes \cite{Bozza2005,Bozza2006,Bozza2007}. This will allow us to widen our study to the Kerr case, where we hope to get similar analytical results that can be useful to get a deeper and thorough comprehension of the contributions of higher order images to the relativistic lines. The existence of wide extended caustics \cite{Bozza2008} hints at an even higher relevance of the contributions of higher order images.
16
7
1607.05365
The shapes of relativistic iron lines observed in spectra of candidate black holes carry the signatures of the strong gravitational fields in which the accretion disks lie. These lines result from the sum of the contributions of all images of the disk created by gravitational lensing, with the direct and first-order images largely dominating the overall shapes. Higher order images created by photons tightly winding around the black holes are often neglected in the modeling of these lines, since they require a substantially higher computational effort. With the help of the strong deflection limit, we present the most accurate semi-analytical calculation of these higher order contributions to the iron lines for Schwarzschild black holes. We show that two regimes exist depending on the inclination of the disk with respect to the line of sight. Many useful analytical formulae can be also derived in this framework.
false
[ "Schwarzschild black holes", "relativistic iron lines", "Higher order images", "gravitational lensing", "Many useful analytical formulae", "these higher order contributions", "the black holes", "the iron lines", "the strong gravitational fields", "sight", "first", "candidate", "Schwarzschild", "respect", "photons", "the accretion disks", "These lines", "the line", "these lines", "the most accurate semi-analytical calculation" ]
15.995081
8.206721
118
12517385
[ "Plucinsky, Paul P.", "Beardmore, Andrew P.", "Foster, Adam", "Haberl, Frank", "Miller, Eric D.", "Pollock, Andrew M. T.", "Sembay, Steve" ]
2017A&A...597A..35P
[ "SNR 1E 0102.2-7219 as an X-ray calibration standard in the 0.5-1.0 keV bandpass and its application to the CCD instruments aboard Chandra, Suzaku, Swift and XMM-Newton" ]
45
[ "Harvard-Smithsonian Center for Astrophysics, MS-3, 60 Garden Street, Cambridge, MA, 02138, USA", "Department of Physics and Astronomy, University of Leicester, Leicester, LE1 7RH, UK", "Harvard-Smithsonian Center for Astrophysics, MS-3, 60 Garden Street, Cambridge, MA, 02138, USA", "Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstraße, 85748, Garching, Germany", "MIT Kavli Institute for Astrophysics and Space Research, Cambridge, MA, 02139, USA", "University of Sheffield, Department of Physics and Astronomy, Hounsfield Road, Sheffield, S3 7RH, UK", "Department of Physics and Astronomy, University of Leicester, Leicester, LE1 7RH, UK" ]
[ "2017A&A...602A..81C", "2017A&A...603A..58R", "2017AN....338..140W", "2017ApJ...840..112S", "2017ApJ...851L..28P", "2017JApA...38...29S", "2018ApJ...868...47K", "2018NatAs...2..465V", "2018PASJ...70...12H", "2018SPIE10699E..6BP", "2018SSRv..214...28R", "2018SSRv..214...48M", "2019A&A...624A..84P", "2019A&A...631A.127M", "2019ApJ...871...64A", "2019ApJ...873...53A", "2019ApJ...874...14X", "2019supe.book..225R", "2020ApJ...896...39A", "2020ApJ...904...70L", "2020MNRAS.491.1585H", "2020PhRvL.125x3001L", "2021AJ....162..254M", "2021ApJ...912...95R", "2021ApJ...916...76A", "2021ApJ...918L..27R", "2021ApJ...923..191P", "2021arXiv211101613M", "2022A&A...661A..25H", "2022A&A...661A..33M", "2022AJ....163..130R", "2022ApJ...924..119W", "2022ApJ...935L..11L", "2022ApJ...941..150S", "2022EPJD...76...38S", "2022JATIS...8d4002G", "2022hxga.book...83S", "2023A&A...676A.142N", "2023ApJ...954..112S", "2023ApJ...959...13R", "2023arXiv230213775K", "2023arXiv230810936S", "2024A&A...682A..34M", "2024arXiv240313933D", "2024arXiv240614466S" ]
[ "astronomy" ]
15
[ "instrumentation: detectors", "X-rays: individuals: 1E 0102.2-7219", "ISM: supernova remnants", "supernovae: general", "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1979ApJ...228..939C", "1981ApJ...243..736S", "1981ApJ...248L.105D", "1983ApJ...268L..11T", "1989ApJ...338..812B", "1989ApJ...342.1207N", "1989ApJS...70..865R", "1990ApJS...74...93R", "1992ApJ...384..508R", "1998SPIE.3444..210B", "2000ApJ...534L..47G", "2000ApJ...537..667B", "2000ApJ...539L..41C", "2000ApJ...542..914W", "2000ApJ...543L..61H", "2000SPIE.4012....2W", "2001A&A...365L...1J", "2001A&A...365L...7D", "2001A&A...365L..18S", "2001A&A...365L..27T", "2001A&A...365L.231R", "2001A&A...365L.237S", "2001A&A...365L.242B", "2001ApJ...556L..91S", "2002ApJ...567..716R", "2002PASP..114....1W", "2003SPIE.4851...28G", "2003SPIE.4851...89P", "2004ApJ...605..230F", "2004ApJ...611.1005G", "2004SPIE.5165..217H", "2004SPIE.5165..497M", "2004SPIE.5501..328D", "2005PASP..117.1144C", "2005SSRv..120..165B", "2006ApJ...641..919F", "2006ApJ...642..260S", "2007AJ....133..147P", "2007ApJ...671L..45B", "2007PASJ...59S..23K", "2008SPIE.7011E..2EP", "2009A&A...494..775G", "2009A&A...496..879M", "2009ASPC..411..234D", "2009ApJ...693..822H", "2009ApJ...697..535P", "2009ApJ...700..579R", "2009MNRAS.397.1177E", "2009PASJ...61S...1O", "2010A&A...523A..22N", "2010ApJ...721..597V", "2011A&A...525A..25T", "2011A&A...534A..20P", "2011AJ....141..129H", "2011ApJ...729L..28P", "2011PASJ...63S.657I", "2011xru..conf..283S", "2012A&A...541A..66S", "2012ApJ...756..128F", "2012SPIE.8443E..12P", "2013A&A...552A..47K", "2014A&A...564A..75R", "2015A&A...573A.128D", "2015A&A...575A..30S" ]
[ "10.1051/0004-6361/201628824", "10.48550/arXiv.1607.03069" ]
1607
1607.03069_arXiv.txt
This paper reports the progress of a working group within the {\em International Astronomical Consortium for High Energy Calibration} (IACHEC) to develop a calibration standard for X-ray astronomy in the bandpass from 0.3 to 1.5 keV. An introduction to the IACHEC organization, its objectives and meetings, may be found at the web page {\tt http://web.mit.edu/iachec/}. Our working group was tasked with selecting celestial sources with line-rich spectra in the 0.3--1.5~keV bandpass which would be suitable cross-calibration targets for the current generation of X-ray observatories. The desire for strong lines in this bandpass stems from the fact that the quantum efficiency and spectral resolution of the current CCD-based instruments is changing rapidly from 0.3 to 1.5~keV but the on-board calibration sources currently in use typically have strong lines at only two energies, 1.5~keV (Al~K$\alpha$) and 5.9~keV (Mn~K$\alpha$). The only option available to the current generation of flight instruments to calibrate possible time variable responses in this bandpass is to use celestial sources. The missions which have been represented in this work are the {\em Chandra X-ray Observatory} (\chan)~\citep{weiss2000,weiss2002}, the {\em X-ray Multimirror Mission} (\xmmn)~\citep{jansen2001}, the {\em ASTRO-E2 Observatory} ({\em Suzaku}), and the {\em Swift} Gamma-ray Burst Mission (\swift)~\citep{gehrels2004}. Data from the following instruments have been included in this analysis: the {\em High-Energy Transmission Grating} (HETG)~\citep{canizares2005} and the {\em Advanced CCD Imaging Spectrometer} (ACIS)~\citep{bautz98,garmire03,garmire92} on \chan, the {\em Reflection Gratings Spectrometers} (RGS)~\citep{denherder2001}, the {\em European Photon Imaging Camera} (EPIC) {\em Metal-Oxide Semiconductor} (EPIC-MOS)~\citep{turner2001} CCDs and the EPIC p-n junction (EPIC-pn)~\citep{strueder2001} CCDs on \xmmn, the {\em X-ray Imaging Spectrometer} (XIS) on {\em Suzaku}, and the {\em X-ray Telescope} (XRT)~\citep{burrows2005,godet2007} on {\em Swift}. Ideal calibration targets would need to possess the following qualities. The source would need to be constant in time, to have a simple spectrum defined by a few bright lines with a minimum of line-blending, and to be extended so that ``pileup'' effects in the CCDs are minimized but not so extended that the off-axis response of the telescope dominates the uncertainties in the response. Our working group focused on supernova remnants (SNRs) with thermal spectra and without a central source such as a pulsar, as the class of source which had the greatest likelihood of satisfying these criteria. We narrowed our list to the Galactic SNR Cas~A, the Large Magellanic Cloud remnant N132D and the Small Magellanic Cloud remnant \name~(hereafter E0102). We discarded Cas~A since it is relatively young ($\sim350$~yr), with significant brightness fluctuations in the X-ray, radio, and optical over the past three decades \citep{patnaude2007,patnaude2009,patnaude2011}, it contains a faint (but apparently variable) central source, and it is relatively large (radius $\sim3.5$~arcminutes). We discarded N132D because it has a complicated, irregular morphology in X-rays \citep{borkowski2007} and its spectrum shows strong, complex Fe emission \citep{behar2001}. The spectrum of N132D is significantly more complicated in the 0.5--1.0~keV bandpass than the spectrum of E0102. We therefore settled on E0102 as the most suitable source given its relatively uniform morphology, small size (radius $\sim0.4$~arcminutes), and comparatively simple X-ray spectrum. We presented preliminary results from this effort in \cite{plucinsky2008} and \cite{plucinsky2012} using a few observations with the calibrations available at that time. In this paper, we present an updated analysis of the representative data acquired early in the various missions and expand our investigations to include a characterization of the time dependence of the response of the various CCD instruments. The low energy responses of some of the instruments (ACIS-S3, EPIC-MOS, \& XIS) included in this analysis have a complicated time dependence due to the time-variable accumulation of a contamination layer. A primary objective of this paper is to inform the Guest Observer communities of the respective missions on the current accuracy of the calibration at these low energies. \medskip
16
7
1607.03069
Context. The flight calibration of the spectral response of charge-coupled device (CCD) instruments below 1.5 keV is difficult in general because of the lack of strong lines in the on-board calibration sources typically available. This calibration is also a function of time due to the effects of radiation damage on the CCDs and/or the accumulation of a contamination layer on the filters or CCDs. <BR /> Aims: We desire a simple comparison of the absolute effective areas of the current generation of CCD instruments onboard the following observatories: Chandra ACIS-S3, XMM-Newton (EPIC-MOS and EPIC-pn), Suzaku XIS, and Swift XRT and a straightforward comparison of the time-dependent response of these instruments across their respective mission lifetimes. <BR /> Methods: We have been using 1E 0102.2-7219, the brightest supernova remnant in the Small Magellanic Cloud, to evaluate and modify the response models of these instruments. 1E 0102.2-7219 has strong lines of O, Ne, and Mg below 1.5 keV and little or no Fe emission to complicate the spectrum. The spectrum of 1E 0102.2-7219 has been well-characterized using the RGS gratings instrument on XMM-Newton and the HETG gratings instrument on Chandra. As part of the activities of the International Astronomical Consortium for High Energy Calibration (IACHEC), we have developed a standard spectral model for 1E 0102.2-7219 and fit this model to the spectra extracted from the CCD instruments. The model is empirical in that it includes Gaussians for the identified lines, an absorption component in the Galaxy, another absorption component in the SMC, and two thermal continuum components with different temperatures. In our fits, the model is highly constrained in that only the normalizations of the four brightest lines/line complexes (the O vii Heα triplet, O viii Lyα line, the Ne ix Heα triplet, and the Ne x Lyα line) and an overall normalization are allowed to vary, while all other components are fixed. We adopted this approach to provide a straightforward comparison of the measured line fluxes at these four energies. We have examined these measured line fluxes as a function of time for each instrument after applying the most recent calibrations that account for the time-dependent response of each instrument. <BR /> Results: We performed our effective area comparison with representative, early mission data when the radiation damage and contamination layers were at a minimum, except for the XMM-Newton EPIC-pn instrument which is stable in time. We found that the measured fluxes of the O vii Heαr line, the O viii Lyα line, the Ne ix Heαr line, and the Ne x Lyα line generally agree to within ±10% for all instruments, with 38 of our 48 fitted normalizations within ± 10% of the IACHEC model value. We then fit all available observations of 1E 0102.2-7219 for the CCD instruments close to the on-axis position to characterize the time dependence in the 0.5-1.0 keV band. We present the measured line normalizations as a function of time for each CCD instrument so that the users may estimate the uncertainty in their measured line fluxes for the epoch of their observations.
false
[ "CCD instruments", "instrument", "strong lines", "time", "contamination layers", "the Ne x Lyα line", "CCD", "EPIC", "each CCD instrument", "the CCD instruments", "the Ne ix Heαr line", "the O viii Lyα line", "radiation damage", "the measured line normalizations", "CCDs", "O viii Lyα line", "the measured line fluxes", "their measured line fluxes", "1E", "different temperatures" ]
9.877862
3.664184
67
3897689
[ "Kulkarni, Girish", "Choudhury, Tirthankar Roy", "Puchwein, Ewald", "Haehnelt, Martin G." ]
2016MNRAS.463.2583K
[ "Models of the cosmological 21 cm signal from the epoch of reionization calibrated with Ly α and CMB data" ]
26
[ "Institute of Astronomy and Kavli Institute of Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "National Centre for Radio Astrophysics, Tata Institute of Fundamental Research, Pune 411007, India", "Institute of Astronomy and Kavli Institute of Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy and Kavli Institute of Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK" ]
[ "2017ApJ...844..117X", "2017ApJ...848...23K", "2017JCAP...11..004A", "2017MNRAS.466.2302B", "2017MNRAS.466.4826K", "2017MNRAS.468..122H", "2018ApJ...865..130V", "2018MNRAS.473..765C", "2018MNRAS.478.1065C", "2018MNRAS.479.2564W", "2018arXiv180608687R", "2019MNRAS.483.2907J", "2019MNRAS.485.1350W", "2019MNRAS.485.3486D", "2019MNRAS.486.4075O", "2020MNRAS.493..855Z", "2020MNRAS.494..703W", "2021FrASS...8...12R", "2021JCAP...01..003R", "2022MNRAS.511.4005K", "2022MNRAS.512..792C", "2022MNRAS.516.3389L", "2022RAA....22c5027S", "2023MNRAS.523.3503P", "2023arXiv231106344S", "2024arXiv240406548A" ]
[ "astronomy" ]
15
[ "galaxies: high-redshift", "intergalactic medium", "dark ages", "reionization", "first stars", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stw2168", "10.48550/arXiv.1607.03891" ]
1607
1607.03891_arXiv.txt
Finally unraveling the complete ionization history of hydrogen with high-redshift 21~cm observations is the major science driver of currently operating and planned low-frequency radio telescopes. Achieving the necessary dynamic range for accurate models of the reionization process has thereby been recognized as a key challenge \citep{2011ASL.....4..228T}. Numerical simulations that aim to be self-consistent in their modeling of ionizing sources and the radiative transfer of ionizing photons in the intergalactic medium (IGM) cannot account for the clustering of sources or the structure of the ionization field on scales greater than $\sim 10$ comoving Mpc \citep{2012MNRAS.427.2464F, 2013MNRAS.436.1818F, 2014ApJ...789..149S, 2015MNRAS.451.1586P}. On the other hand, simulations that can take these large scale effects into consideration have low spatial and mass resolution and are unable to consistently model small-scale effects such as radiative feedback on ionizing sources and self-shielding of high density regions in the IGM \citep{2007ApJ...659..865L, 2009MNRAS.393...32T, 2010ApJ...724..244A, 2012ApJ...756L..16A, 2012AIPC.1480..248S, 2015MNRAS.453.3593B, 2015MNRAS.454.1012A}. Simulations are nevertheless crucial for the ongoing and upcoming experiments that aim to detect the 21~cm signal from the epoch of reionization \citep{2015aska.confE...7I}. The 21~cm brightness distribution is expected to eventually become the ultimate probe of reionization. The design of instruments capable of detecting this signal is guided by its predictions from numerical simulations \citep[e.g.,][]{2012ApJ...753...81P}. Simulations are also crucial in the interpretation of the results of these experiments \citep{2016MNRAS.455.4295G}, all of which aim to detect the 21~cm signal statistically. Thus, given their relevance, not only should these simulations have a large enough dynamic range to be self-consistent and convergent but they should also be consistent with other currently available constraints on the epoch of reionization. Due to their computational cost, most simulations of the 21~cm signal lack one or both of these properties. It has been argued that the goal of self-consistent large scale simulation of cosmic reionization is now gradually coming within reach thanks to Moore's Law \citep{2014ApJ...793...29G, 2014ApJ...793...30G, 2015arXiv151100011O, 2015ApJS..216...16N}, but semi-numerical and analytical methods of reionization modeling continue to remain attractive for efficient and flexible exploration of the parameter space, especially given the paucity of data at high redshifts \citep{2007ApJ...669..663M, 2008MNRAS.386.1683G, 2009MNRAS.394..960C, 2011MNRAS.411..955M, 2011MNRAS.412.2781K, 2011MNRAS.417.2264V, 2012MNRAS.423..862K, 2012ApJ...747..126A, 2013MNRAS.428L...1M, 2013RAA....13..373Z, 2013ApJ...776...81B, 2013ApJ...768...71R, 2013ApJ...771...35K, 2014MNRAS.440.1662S, 2014MNRAS.442.1470P, 2015MNRAS.454L..76M, 2016MNRAS.457.1550H}. In this paper, we combine a high dynamic range cosmological simulation with an excursion set based model for the growth of ionized regions to predict the 21~cm signal during the epoch of reionization. We follow the approach of \citeauthor{2015MNRAS.452..261C} (\citeyear{2015MNRAS.452..261C}; hereafter CPHB15) to calibrate the simulation parameters such that they reproduce the IGM Lyman-$\alpha$ (Ly$\alpha$) opacity at $z \lesssim 6$ \citep{2006ARA&A..44..415F, 2015MNRAS.447.3402B, 2015MNRAS.447..499M} as well as the cosmic microwave background (CMB) constraints on the electron scattering optical depth \citep{2016arXiv160502985P, 2016arXiv160503507P}. The advantage of this approach is that once the reionization history is given, all other quantities of interest---such as the photoionization rate, emissivity of ionizing sources, and the clumping factor of the IGM---can be calculated self-consistently from the simulation box at each redshift. This enables us to simulate concordant models of reionization consistent with a wide variety of observations without losing the dynamic range of our simulations. CPHB15 applied this method to study the evolution of Ly$\alpha$ emission in high-redshift galaxies by calibrating a ``hybrid'' cosmological simulation box, which was created by combining a low-resolution cosmological simulation box at large scales with a high-resolution simulation box at small scales. A similar approach was used by \citet{2015MNRAS.446..566M} to study the evolution of the Lyman-$\alpha$ emitter fraction of high-redshift galaxies. In their approach, a seminumerical scheme was used to obtain the low-resolution simulation. The hybrid box used in CPHB15 formally had very high dynamic range (equivalent to a cosmological simulation with $2\times 5120^3$ particles in a 100 $h^{-1}$cMpc box) but did not correctly represent the clustering of matter at scales larger than the size of the small box, which was 10 $h^{-1}$cMpc. Thus from the point of view of deriving the cosmological 21~cm signal, hybrid boxes are of little use as they fail to yield, e.g., the 21~cm power spectrum at scales of interest. The main improvement in the simulation method used in this work is the use of a cosmological hydrodynamical simulation with improved dynamic range. We follow CPHB15 and consider three different reionization histories for our analysis. One of our reionization histories follows the widely used model of the meta-galactic UV background by \citeauthor{2012ApJ...746..125H} (\citeyear{2012ApJ...746..125H}; hereafter HM12). This model was tuned to match the constraints on reionization from the Wilkinson Microwave Anisotropy Probe (WMAP; \citealt{2011ApJS..192...14J}) and predicts an electron scattering optical depth higher than the recent Planck measurements \citep{2016arXiv160502985P, 2016arXiv160503507P}. We therefore explore two other reionization histories in which reionization is completed later than in the HM12 model and the electron scattering predictions are consistent with Planck results. In addition to the evolution of the ionized fraction, the 21~cm signal also depends on the distribution of optically thick systems that are self-shielded from the ionizing radiation. Such systems, which are high-redshift counterparts of the Lyman-limit systems seen in quasar absorption spectra, are usually missed by low resolution simulations. We leverage our high dynamic range to study the effect of these self-shielded regions on the 21~cm signal using a prescription for self-shielding provided by \citet{2013MNRAS.430.2427R}. Finally, we consider whether our predicted signal can be observed by five ongoing and upcoming 21~cm experiments. The main aim of the paper is thus to use models that are calibrated to existing data and predict the 21~cm signal and its detectability at different redshifts. We describe our simulations and the calibration procedure in Section~\ref{sec:sims}. Section~\ref{sec:21cm} presents our predictions for the 21~cm signal and its observability in ongoing and future experiments. We investigate the effect of various assumptions on our results in Section~\ref{sec:ss_effect} and conclude with a discussion in Section~\ref{sec:end}. Our $\Lambda$CDM cosmological model has $\Omega_\mathrm{b}=0.0482$, $\Omega_\mathrm{m}=0.308$, $\Omega_\Lambda=0.692$, $h=0.678$, $n=0.961$, $\sigma_8=0.829$, and $Y_\mathrm{He}=0.24$ \citep{2014A&A...571A..16P}. \begin{figure*} \begin{center} \begin{tabular}{cc} \includegraphics*[width=\columnwidth]{{halos_threshold2.0}.pdf} & \includegraphics*[width=\columnwidth]{{halos_threshold0.997}.pdf} \end{tabular} \end{center} \caption{Distribution of gas density at $z=7$. Black symbols denote the locations of centres of masses of dark matter haloes. Shaded areas in both panels show ionized regions identified by the excursion set method. Colour scale shows the gas density. The left panel shows the ionization field for $\zeta_\mathrm{eff}=0.5$, which corresponds to $Q_V=0.31$. The right panel shows the ionization field for $\zeta_\mathrm{eff}=1.0$, which corresponds to $Q_V=0.94$.} \label{fig:slices} \end{figure*}
\label{sec:end} We have presented here predictions of the spatial distribution of the 21~cm brightness temperature from the epoch of reionization based on high dynamic range cosmological hydrodynamical simulations from the Sherwood simulation suite \citep{2016arXiv160503462B} for reionization histories motivated by constraints from \lya absorption and emission data as well as CMB data. Our models of the 21~cm signal were obtained by combining the high dynamic range cosmological simulations with excursion set based models of the growth of ionized regions during reionization. This has allowed us to efficiently obtain realistic 21~cm predictions that are firmly anchored in current constraints on how reionization ends and that have sufficient resolution to account for the self-shielding of neutral hydrogen in dense regions within otherwise fully ionized regions. Our main conclusions are as follows: \begin{itemize} \item Our preferred `Late/Default' model of reionization is consistent with a variety of observational constraints such as the electron scattering optical depth, Ly$\alpha$ opacity, galaxy UV luminosity function at high redshifts, and estimates of the hydrogen photoionization rates from quasar absorption spectra \citep{2015MNRAS.453.2943C}. In this model the volume-weighted ionization fraction $Q_V$ evolves from 0.37 at $z=10$ to 0.82 at $z=7$ and reionization is complete at redshift $z\sim 6$. The variance of the 21~cm brightness temperature field at large scales accessible to current and upcoming experiments ($k\sim 0.1$ cMpc$^{-1}h$) reduces from about 20 mK$^2$ at $z=10$ to about 10 mK$^2$ at $z=7$. Power at these large scales first increases from $z=10$ to $8$, when $Q_V$ is 0.65, and then decreases down to $z=7$. The change from the rise to the fall of the signal of reionization, where the 21~cm power peaks occurs when $Q_V$ is about 50\%, which occurs at about $z=8.5$ in this model. The small scale 21~cm power in this model decreases continuously from about 50 mK$^2$ at $z=10$ to 10 mK$^2$ at $z=7$ ($k\sim 1$ cMpc$^{-1}h$). These small scales are, however, unlikely to be accessible to any of the five experiments that we have investigated here, particularly at $z=10$. \item Self-shielding in high density regions within ionized regions affects the 21~cm power in two ways. At the large scales accessible to experiments, self-shielding decreases power by up to 30\%. The contribution to the 21~cm power from self-shielded regions tends to be greater at high redshifts unless the ionization fractions are too small. At small scales, self-shielding enhances the 21~cm power by adding structure within ionized regions. This effect is generally stronger at lower redshift due to the larger volume occupied by ionized regions. The enhancement in power at small scales due to self-shielding can be considerable, often even greater than an order of magnitude. Unfortunately, these scales are too small to be detected by the 21~cm experiments considered here. By suppressing power at intermediate scales self-shielding can reduce the rise and fall signature of the epoch of reionization. \item In addition to our favoured Late/Default model, we have considered two other reionization histories. In the reionization history predicted by the widely used HM12 model of the meta-galactic UV background, reionization occurs earlier and the 21~cm signal peaks as early as $z=10$. Otherwise the evolution of the 21~cm brightness distribution is qualitatively similar to the Late/Default model on all scales. As a result of the earlier rise of the ionized volume fraction the large scale power does not show the same rise and fall behaviour in this model between $z=10$ and $7$ due to the relatively larger ionization fraction already at $z=10$. Instead the large scale variance of the 21~cm brightness decreases continuously from about 30 mK$^2$ at $z=10$ to about 3 mK$^2$ at $z=7$ in this model. Our third reionization history, the `Very Late' model, in which reionization also ends at $z=6$ but starts later than in the other models, also does not show the rise and fall behaviour in the large scale 21~cm power. In this model the large scale power increases continuously from about 2 mK$^2$ at $z=10$ to about 15 mK$^2$ at $z=7$. The effect of self-shielding on the power spectrum in these two models is qualitatively similar to that in the Late model. In all three cases the effect of self-shielding reduces the large scale 21~cm power by up to 30\% and increases the small scale power by factors of up to 10. The latter effect is not accessible at the resolution of the 21~cm experiments we have considered here. \item We have also varied the scaling of the luminosity of the ionizing sources with host halo mass in our models. Our `Nonlinear' model has the same reionization history as the Late/Default model but an ionizing photon rate that is assumed to be a nonlinear function of the halo mass, as, e.g., preferred by the radiation hydrodynamical simulations of \citet{2011MNRAS.410.1703F}. This moderately strengthens the relative role of high mass haloes compared to low mass haloes during reionization. The results of this model are nearly identical to those from our Late/Default model. We have further considered a higher cut-off in the mass of haloes hosting ionizing sources. In this `High Mass' model, we remove low-mass haloes from our ionization field calculation. This scenario corresponds to strong feedback, completely suppressing star formation in low-mass haloes. We find that although this model is calibrated to the same reionization history as our Late/Default model, suppression of low luminosity sources changes the 21~cm signal significantly. Broad features in the power spectrum are qualitatively similar to those in the Late/Default model, but at redshift $z=10$ the large scale power is enhanced almost ten times to about $100$ mK$^2$. Thus reionization histories with strongly clustered bright sources would clearly favour the detection of the 21~cm signal. \item We have derived sensitivity limits of five current and upcoming experiments, PAPER, MWA, LOFAR, HERA, and SKA1-LOW, for the 21~cm power spectrum due to thermal noise and consider the prospects of detecting the reduced large scale 21~cm power spectrum in presence of self-shielding. Assuming perfect foreground removal, LOFAR, HERA, and SKA1-LOW should be able to detect the power spectrum at $k\sim 0.1$ cMpc$^{-1}h$ at $z=8$ at 20-$\sigma$ to 100-$\sigma$ significance in our Late/Default model, assuming perfect foreground removal. The significance drops at $z=7$ as well as $z=10$ due to changes in the ionization structure and evolution in experimental sensitivities. Detection is more difficult with the first generation experiments PAPER and MWA except at scales larger than our box size of 160 $h^{-1}$cMpc, although foreground removal can be difficult at these large scales \citep{2014ApJ...782...66P}. At $k=0.1$ cMpc$^{-1}h$, MWA should be able to detect the rise and fall signal of the epoch of reionization at less than 2-$\sigma$, whereas LOFAR should be able to detect it at nearly 10-$\sigma$. Assuming ideal foreground removal, HERA and SKA1-LOW should be able to detect this signature comfortably at excess of 50-$\sigma$. At their design sensitivity LOFAR, HERA and SKA1-LOW should therefore easily discriminate between the ionization histories presented here. \end{itemize} The calibration procedure used in this paper provides a relatively low-cost method of performing high dynamic range simulations of the cosmological 21~cm signal for reionization histories that are well anchored in constraints from other data on how reionization ends. Although self-consistent large scale simulations of cosmic reionization are now gradually becoming possible, the method presented in this paper is valuable for an efficient and flexible exploration of the relevant parameter space that will be necessary for inference from the statistical detection of the 21~cm signal.
16
7
1607.03891
We present here 21 cm predictions from high dynamic range simulations for a range of reionization histories that have been tested against available Ly α and cosmic microwave background (CMB) data. We assess the observability of the predicted spatial 21 cm fluctuations by ongoing and upcoming experiments in the late stages of reionization in the limit in which the hydrogen spin temperature is significantly larger than the CMB temperature. Models consistent with the available Ly α data and CMB measurement of the Thomson optical depth predict typical values of 10-20 mK<SUP>2</SUP> for the variance of the 21 cm brightness temperature at redshifts z = 7-10 at scales accessible to ongoing and upcoming experiments (k ≲ 1 cMpc<SUP>-1</SUP>h). This is within a factor of a few magnitude of the sensitivity claimed to have been already reached by ongoing experiments in the signal rms value. Our different models for the reionization history make markedly different predictions for the redshift evolution and thus frequency dependence of the 21 cm power spectrum and should be easily discernible by Low-Frequency Array (and later Hydrogen Epoch of Reionization Array and Square Kilometre Array1) at their design sensitivity. Our simulations have sufficient resolution to assess the effect of high-density Lyman limit systems that can self-shield against ionizing radiation and stay 21 cm bright even if the hydrogen in their surroundings is highly ionized. Our simulations predict that including the effect of the self-shielded gas in highly ionized regions reduces the large-scale 21 cm power by about 30 per cent.
false
[ "ongoing experiments", "typical values", "ionizing radiation", "CMB measurement", "Reionization Array", "redshifts z", "CMB", "high dynamic range simulations", "reionization histories", "Square Kilometre", "scales", "ongoing and upcoming experiments", "cMpc", "reionization", "the hydrogen spin temperature", "the CMB temperature", "SUP>2</SUP", "21 cm predictions", "Thomson", "the signal rms value" ]
13.17124
3.547372
120
4384367
[ "Yang, C.", "Omont, A.", "Beelen, A.", "González-Alfonso, E.", "Neri, R.", "Gao, Y.", "van der Werf, P.", "Weiß, A.", "Gavazzi, R.", "Falstad, N.", "Baker, A. J.", "Bussmann, R. S.", "Cooray, A.", "Cox, P.", "Dannerbauer, H.", "Dye, S.", "Guélin, M.", "Ivison, R.", "Krips, M.", "Lehnert, M.", "Michałowski, M. J.", "Riechers, D. A.", "Spaans, M.", "Valiante, E." ]
2016A&A...595A..80Y
[ "Submillimeter H<SUB>2</SUB>O and H<SUB>2</SUB>O<SUP>+</SUP>emission in lensed ultra- and hyper-luminous infrared galaxies at z 2-4" ]
57
[ "Purple Mountain Observatory/Key Lab of Radio Astronomy, Chinese Academy of Sciences, 210008, Nanjing, PR China ; Institut d'Astrophysique Spatiale, CNRS, Univ. Paris-Sud, Université Paris-Saclay, Bât. 121, 91405, Orsay Cedex, France; Graduate University of the Chinese Academy of Sciences, 19A Yuquan Road, Shijingshan District, 10049, Beijing, PR China; CNRS, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France; UPMC Univ. Paris 06, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France", "CNRS, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France; UPMC Univ. Paris 06, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France", "Institut d'Astrophysique Spatiale, CNRS, Univ. Paris-Sud, Université Paris-Saclay, Bât. 121, 91405, Orsay Cedex, France", "Universidad de Alcalá, Departamento de Fisica y Matemáticas, Campus Universitario, 28871 Alcalá de Henares, Madrid, Spain", "Institut de Radioastronomie Millimétrique (IRAM), 300 rue de la Piscine, 38406 Saint-Martin-d', Hères, France", "Purple Mountain Observatory/Key Lab of Radio Astronomy, Chinese Academy of Sciences, 210008, Nanjing, PR China", "Leiden Observatory, Leiden University, PO Box 9513, 2300 RA, Leiden, The Netherlands", "Max Planck Institut für Radioastronomie, Auf dem Hgel 69, 53121, Bonn, Germany", "CNRS, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France; UPMC Univ. Paris 06, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, 43992, Onsala, Sweden", "Department of Physics and Astronomy, Rutgers, The State University of New Jersey, 136 Frelinghuysen Road, Piscataway, NJ, 08854-8019, USA", "Astronomy Department, Cornell University, 220 Space Sciences Building, Ithaca, NY, 14853, USA", "Department of Physics and Astronomy, University of California, Irvine, Irvine, CA, 92697, USA", "Joint ALMA Observatory, 3107 Alonso de Córdova, Vitacura, Santiago, Chile", "Universitat Wien, Institut für Astrophysik, Türkenschanzstrasse 17, 1180, Wien, Austria", "School of Physics and Astronomy, University of Nottingham, University Park, Nottingham, NG7 2RD, UK", "Institut de Radioastronomie Millimétrique (IRAM), 300 rue de la Piscine, 38406 Saint-Martin-d', Hères, France", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh, EH9 3HJ, UK; European Southern Observatory, Karl Schwarzschild Straße 2, 85748, Garching, Germany", "Institut de Radioastronomie Millimétrique (IRAM), 300 rue de la Piscine, 38406 Saint-Martin-d', Hères, France", "CNRS, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France; UPMC Univ. Paris 06, UMR 7095, Institut d'Astrophysique de Paris, 75014, Paris, France", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh, EH9 3HJ, UK", "Astronomy Department, Cornell University, 220 Space Sciences Building, Ithaca, NY, 14853, USA", "Kapteyn Astronomical Institute, University of Groningen, 9700 AV, Groningen, The Netherlands", "School of Physics and Astronomy, Cardiff University, The Parade, Cardiff, CF24 3AA, UK" ]
[ "2017A&A...608A.144Y", "2017ApJ...837...12W", "2017ApJ...846....5L", "2017ApJ...850....1R", "2017ApJ...850..170O", "2017MNRAS.465.3558N", "2017Natur.548..430F", "2017PASA...34...54G", "2017PhDT........21Y", "2017PhDT.......174M", "2018A&A...615A.142A", "2018ApJ...856..174V", "2018ApJ...864L..11X", "2018ApJ...865..127I", "2018MNRAS.475.3467E", "2018MNRAS.481...59Z", "2018NatAs...2...56Z", "2019A&A...624A.138Y", "2019A&A...628A..23A", "2019ApJ...880...92J", "2019ApJ...880..153Y", "2019ApJ...887..144J", "2019sf2a.conf....3B", "2020A&A...634L...3Y", "2020A&A...641A.124L", "2020A&A...642A.155Z", "2020ApJ...889..162L", "2020ApJ...898...33C", "2020IAUS..352..297Y", "2020RSOS....700556H", "2021A&A...645A..45C", "2021A&A...645A..49G", "2021A&A...646A.122B", "2021A&A...646A.178S", "2021A&A...648A..24V", "2021A&A...652A..66P", "2021ApJ...921...97J", "2022A&A...665A.107T", "2022A&A...666L...3G", "2022A&A...667A...9P", "2022A&A...667A.135D", "2022FrASS...987567G", "2022MNRAS.510.3734D", "2023A&A...673A.157D", "2023A&A...674A..95D", "2023A&A...678A..28B", "2023A&A...680A..95Y", "2023ApJ...948...44R", "2023ApJ...951...32T", "2023ApJ...952...90P", "2023MNRAS.521.5508H", "2024A&A...684A..56K", "2024ApJ...962..119L", "2024Galax..12...18P", "2024MNRAS.527.6321Q", "2024MNRAS.528.6222B", "2024arXiv240105487R" ]
[ "astronomy" ]
14
[ "galaxies: high-redshift", "galaxies: ISM", "infrared: galaxies", "submillimeter: galaxies", "radio lines: ISM", "ISM: molecules", "Astrophysics - Astrophysics of Galaxies" ]
[ "1992ARA&A..30..575C", "1992ApJ...387L..55S", "1995ApJ...450..559B", "1996A&A...315L..27K", "1998ARA&A..36..189K", "2001ApJ...554..803Y", "2003ARA&A..41..241D", "2003ApJS..148..175S", "2004ApJ...606..271G", "2004ApJ...613..247G", "2005A&A...436..397M", "2007ApJ...660L..93G", "2007ApJ...665.1489K", "2007MNRAS.378..983B", "2008ApJ...675..303G", "2009ASPC..411..251M", "2010A&A...518L...1P", "2010A&A...518L..42V", "2010A&A...518L..43G", "2010A&A...518L.110G", "2010A&A...521L...1W", "2010A&A...521L..10N", "2010AJ....140.1868W", "2010PASP..122..499E", "2010SPIE.7731E..16N", "2010Sci...330..800N", "2011A&A...525A.119M", "2011A&A...530L...3O", "2011ApJ...738L...6L", "2011ApJ...740L..15M", "2011ApJ...741L..37B", "2011ApJ...741L..38V", "2011ApJ...743...94R", "2011MNRAS.415..911P", "2012A&A...538L...4C", "2012A&A...539A...8G", "2012A&A...541A...4G", "2012ApJ...752..152H", "2012ApJ...753...70K", "2012ApJ...757..135L", "2012ApJ...758..108S", "2012MNRAS.424.2429S", "2013A&A...550A..25G", "2013A&A...551A.115O", "2013ApJ...762L..16M", "2013ApJ...767...88W", "2013ApJ...768...55P", "2013ApJ...771L..24Y", "2013ApJ...772..137I", "2013ApJ...779...25B", "2013ApJ...779...67B", "2013ApJ...779L..19P", "2013ChRv..113.9043V", "2013Natur.495..344V", "2013Natur.496..329R", "2014A&A...567A..91G", "2014ApJ...783...59R", "2014ApJ...790...40C", "2014ApJ...797..138C", "2014PhR...541...45C", "2015A&A...575A..56B", "2015ApJ...808L...4A", "2015ApJ...810L..14L", "2015ApJ...814....9K", "2015MNRAS.451L..40R", "2015MNRAS.452.2258D", "2016A&A...585A.103S" ]
[ "10.1051/0004-6361/201628160", "10.48550/arXiv.1607.06220" ]
1607
1607.06220_arXiv.txt
\label{section:Introduction} After molecular hydrogen (H$_2$) and carbon monoxide (CO), the water molecule (H$_2$O) can be one of the most abundant molecules in the interstellar medium (ISM) in galaxies. It provides some important diagnostic tools for various physical and chemical processes in the ISM \citep[e.g.][and references therein]{2013ChRv..113.9043V}. Prior to the {\it Herschel Space Observatory} \citep{2010A&A...518L...1P}, in extragalactic sources, non-maser \hto\ rotational transitions were only detected by the \textit{Infrared Space Observatory} \citep[\textit{ISO},][]{1996A&A...315L..27K} in the form of far-infrared absorption lines \citep{2004ApJ...613..247G, 2008ApJ...675..303G}. Observations of local infrared bright galaxies by {\it Herschel} have revealed a rich spectrum of submillimeter (submm) \hto\ emission lines (submm \hto\ refers to rest-frame submillimeter \hto\ emission throughout this paper if not otherwise specified). Many of these lines are emitted from high-excitation rotational levels with upper-level energies up to $E_\mathrm{up}$/$k = 642\,$K \citep[e.g.][]{2010A&A...518L..42V, 2010A&A...518L..43G, 2012A&A...541A...4G, 2013A&A...550A..25G, 2011ApJ...743...94R, 2012ApJ...753...70K, 2012ApJ...758..108S, 2013ApJ...762L..16M, 2013ApJ...779L..19P, 2013ApJ...768...55P}. Excitation analysis of these lines has revealed that they are probably excited through absorption of \fir\ photons from thermal dust emission in warm dense regions of the ISM \citep[e.g.][]{2010A&A...518L..43G}. Therefore, unlike the canonical CO lines that trace collisional excitation of the molecular gas, these \hto\ lines represent a powerful diagnostic of the \fir\ radiation field. Using the {\it Herschel} archive data, \citet[][hereafter Y13]{2013ApJ...771L..24Y} have undertaken a first systematic study of submm \hto\ emission in local \ir\ galaxies. \hto\ was found to be the strongest molecular emitter after CO within the submm band in those \ir-bright galaxies, even with higher flux density than that of CO in some local ULIRGs (velocity-integrated flux density of \htot321312 is larger than that of \co54 in four galaxies out of 45 in the \citetalias{2013ApJ...771L..24Y} sample). The luminosities of the submm \hto\ lines (\lhto) are near-linearly correlated with total \ir\ luminosity ($L_\mathrm{IR}$, integrated over 8--1000\,$\mu$m) over three orders of magnitude. The correlation is revealed to be a straightforward result of \fir\ pumping: \hto\ molecules are excited to higher energy levels through absorbing \fir\ photons, then the upper level molecules cascade toward the lines we observed in an almost constant fraction (Fig.\,\ref{fig:h2o-e-level}). Although the galaxies dominated by active galactic nuclei (AGN) have somewhat lower ratios of \lhto/\lir, there does not appear to be a link between the presence of an AGN and the submm \hto\ emission \citepalias{2013ApJ...771L..24Y}. The \hto\ emission is likely to trace the \fir\ radiation field generated in star-forming nuclear regions in galaxies, explaining its tight correlation with \fir\ luminosity. Besides detections of the \hto\ lines in local galaxies from space telescopes, redshifted submm \hto\ lines in \hz\ lensed Ultra- and Hyper-Luminous InfraRed Galaxies (ULIRGs, $10^{13}\,L_\odot > L_\mathrm{IR} \geq 10^{12}$\,\lsun; HyLIRGs, $L_\mathrm{IR} \geq 10^{13}$\,\lsun) can also be detected by ground-based telescopes in atmospheric windows with high transmission. Strong gravitational lensing boosts the flux and allows one to detect the \hto\ emission lines easily. Since our first detection of submm \hto\ in a lensed {\it Herschel} source at $z = 2.3$ \citep{2011A&A...530L...3O} using the IRAM NOrthern Extended Millimeter Array (NOEMA), several individual detections at \hz\ have also been reported \citep{2011ApJ...738L...6L, 2011ApJ...741L..38V, 2011ApJ...741L..37B, 2012A&A...538L...4C, 2012ApJ...757..135L, 2013ApJ...779...67B, 2013A&A...551A.115O, 2013Natur.495..344V, 2013ApJ...767...88W, 2014ApJ...783...59R}. These numerous and easy detections of \hto\ in \hz\ lensed ULIRGs show that its lines are the strongest submm molecular lines after CO and may be an important tool for studying these galaxies. We have carried out a series of studies focussing on submm \hto\ emission in \hz\ lensed galaxies since our first detection. Through the detection of $J=2$ \hto\ lines in seven \hz\ lensed Hy/ULIRGs reported by \citet[][hereafter O13]{2013A&A...551A.115O}, a slightly super-linear correlation between \lhto\ and \lir\ (\lhto\;$\propto$\;\lir$^{1.2}$) from local ULIRGs and \hz\ lensed Hy/ULIRGs has been found. This result may imply again that \fir\ pumping is important for \hto\ excitation in \hz\ extreme starbursts. The average ratios of \lhto\ to \lir\ for the $J=2$ \hto\ lines in the \hz\ sources tend to be $1.8\pm0.9$ times higher than those seen locally \citepalias{2013ApJ...771L..24Y}. This shows that the same physics with infrared pumping should dominate \hto\ excitation in ULIRGs at low and high redshift, with some specificity at \hz\ probably linked to the higher luminosities. Modelling provides additional information about the \hto\ excitation. For example, through LVG modelling, \cite{2013Natur.496..329R} argue that the excitation of the submm \hto\ emission in the $z \sim 6.3$ submm galaxy is \fir\ pumping dominated. Modelling of the local {\it Herschel} galaxies of \citetalias{2013ApJ...771L..24Y} has been carried out by \citet[][hereafter G14]{2014A&A...567A..91G}. They confirm that \fir\ pumping is the dominant mechanism responsible for the submm \hto\ emission (except for the ground-state emission transitions, such as para-\hto\ transition \t111000) in the extragalactic sources. Moreover, collisional excitation of the low-lying ($J \leq 2$) \hto\ lines could also enhance the radiative pumping of the ($J \geq 3$) high-lying lines. The ratio between low-lying and high-lying \hto\ lines is sensitive to the dust temperature (\td) and \hto\ column density ($N_\mathrm{H_2O}$). From modelling the average of local star-forming- and mild-AGN-dominated galaxies, \citetalias{2014A&A...567A..91G} show that the submm \hto\ emission comes from regions with $N_\mathrm{H_2O} \sim (0.5\text{--}2) \times 10^{17}$\,cm$^{-2}$ and a 100\,$\mu$m continuum opacity of $\tau_{100} \sim 0.05\text{--}0.2$, where \hto\ is mainly excited by warm dust with a temperature range of $45\text{--}75$\,K. \hto\ lines thus provide key information about the properties of the dense cores of ULIRGs, that is, their \hto\ content, the infrared radiation field and the corresponding temperature of dust that is warmer than the core outer layers and dominates the far-infrared emission. Observations of the submm \hto\ emission, together with appropriate modelling and analysis, therefore allows us to study the properties of the \fir\ radiation sources in great detail. So far, the excitation analysis combining both low- and high-lying \hto\ emission has only been done in a few case studies. Using \hto\ excitation modelling considering both collision and \fir\ pumping, \cite{2010A&A...518L..43G} and \cite{2011ApJ...741L..38V} estimate the sizes of the \fir\ radiation fields in Mrk\,231 and APM\,08279+5255 (APM\,08279 hereafter), which are not resolved by the observations directly, and suggest their AGN dominance based on their total enclosed energies. This again demonstrates that submm \hto\ emission is a powerful diagnostic tool which can even transcend the angular resolution of the telescopes. The detection of submm \hto\ emission in the {\it Herschel}-ATLAS\footnote{The {\it Herschel}-ATLAS is a project with {\it Herschel}, which is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA. The {\it H}-ATLAS website is \url{http://www.h-atlas.org}.} \citep[][{\it H}-ATLAS hereafter]{2010PASP..122..499E} sources through gravitational lensing allows us to characterise the \fir\ radiation field generated by intense star-forming activity, and possibly AGN, and learn the physical conditions in the warm dense gas phase in extreme starbursts in the early Universe. Unlike standard dense gas tracers such as HCN, which is weaker at \hz\ compared to that of local ULIRGs \citep{2007ApJ...660L..93G}, submm \hto\ lines are strong and even comparable to high-$J$ CO lines in some galaxies \citepalias{2013ApJ...771L..24Y, 2013A&A...551A.115O}. Therefore, \hto\ is an efficient tracer of the warm dense gas phase that makes up a major fraction of the total molecular gas mass in \hz\ Hy/ULIRGs \citep{2014PhR...541...45C}. The successful detections of submm \hto\ lines in both local \citepalias{2013ApJ...771L..24Y} and the \hz\ universe \citepalias{2013A&A...551A.115O} show the great potential of a systematic study of \hto\ emission in a large sample of \ir\ galaxies over a wide range in redshift (from local up to $z\sim4$) and luminosity ($\lir \sim10^{10}$--$10^{13}$\,\lsun). However, our previous \hz\ sample was limited to seven sources and to one $J=2$ para-\hto\ line ($E_\mathrm{up}$/$k = 100$--$127\,$K) per source \citepalias{2013A&A...551A.115O}. In order to further constrain the conditions of \hto\ excitation, to confirm the dominant role of \fir\ pumping and to learn the physical conditions of the warm dense gas phase in \hz\ starbursts, it is essential to extend the studies to higher excitation lines. We thus present and discuss here the results of such new observations of a strong $J=3$ ortho-\hto\ line with $E_\mathrm{up}$/$k = 304\,$K in six strongly lensed {\it H}-ATLAS galaxies at z\,$\sim$\,2.8--3.6, where a second lower-excitation $J=2$ para-\hto\ line was also observed (Fig.\,\ref{fig:h2o-e-level} for the transitions and the corresponding $E_\mathrm{up}$). \begin{figure}[htbp] \begin{center} \includegraphics[scale=0.441]{e_level} \caption{ Energy level diagrams of \hto\ and \htop\ shown in black and red, respectively. Dark blue arrows are the submm \hto\ transitions we have observed in this work. Pink dashed lines show the \fir\ pumping path of the \hto\ excitation in the model we use, with the wavelength of the photon labeled. The light blue dashed arrow is the transition from para-\hto\ energy level $2_{20}$ to $2_{11}$ along the cascade path from $2_{20}$ to $1_{11}$. Rotational energy levels of \hto\ and \htop, as well as fine structure component levels of \htop\, are also shown in the figure. } \label{fig:h2o-e-level} \end{center} \end{figure} We describe our sample, observation and data reduction in Section \ref{section:Sample and Observation}. The observed properties of the \hz\ submm \hto\ emission are presented in Section \ref{Results}. Discussions of the lensing properties, \lhto-\lir\ correlation, \hto\ excitation, comparison between \hto\ and CO, AGN contamination will be given in Section \ref{Discussion}. Section \ref{htop} describes the detection of \htop\ lines. We summarise our results in Section \ref{Conclusions}. A flat $\Lambda$CDM cosmology with $H_{0}=71\,{\rm km\,s^{-1}\,Mpc^{-1}}$, $\Omega_{M}=0.27$, $\Omega_{\Lambda}=0.73$ \citep{2003ApJS..148..175S} is adopted throughout this paper.
\normalsize \label{Conclusions} In this paper, we report a survey of submm \hto\ emission at redshift $z \sim 2\text{--}4$, by observing a higher excited ortho-\htot321312 in 6 sources and several complementary $J=2$ para-\hto\ emission lines in the warm dense cores of 11 \hz\ lensed extreme starburst galaxies (Hy/ULIRGs) discovered by {\it H}-ATLAS. So far, we have detected an \hto\ line in most of our observations of a total sample of 17 \hz\ lensed galaxies, in other words, we have detected both $J=2$ para-\hto\ and $J=3$ ortho-\hto\ lines in five, and in ten other sources only one $J=2$ para-\hto\ line. In these \hz\ Hy/ULIRGs, \hto\ is the second strongest molecular emitter after CO within the submm band, as in local ULIRGs. The spatially integrated \hto\ emission lines have a velocity-integrated flux density ranging from 4 to 15\,Jy\,km\,s$^{-1}$, which yields the apparent \hto\ emission luminosity, $\mu$\lhto\, ranging from $\sim 6\text{--}22 \times 10^{8}$\,\lsun. After correction for gravitation lensing magnification, we obtained the intrinsic \lhto\ for para-\hto\ lines \t202111, \t211202 and ortho-\htot321312. The luminosities of the three \hto\ lines increase with \lir\ as \lhto\;$\propto$\;\lir$^{1.1\text{--}1.2}$. This correlation indicates the importance of \fir\ pumping as a dominant mechanism of \hto\ excitation. Comparing with $J=3$ to $J=6$ CO lines, the linewidths between \hto\ and CO are similar, and the velocity-integrated flux densities of \hto\ and CO are comparable. The similarity in line profiles suggests that these two molecular species possibly trace similar intense star-forming regions. Using the \fir\ pumping model, we have analysed the ratios between $J=2$ and $J=3$ \hto\ lines and \lhto/\lir\ in 5 sources with both $J$ \hto\ lines detected. We have derived the ranges of the warm dust temperature ($T_\mathrm{warm}$), the \hto\ column density per unit velocity interval (\nhto/$\Delta V$) and the optical depth at 100\,$\mu$m ($\tau_{100}$). Although there are strong degeneracies, these modelling efforts confirm that, similar to those of local ULIRGs, these submm \hto\ emissions in \hz\ Hy/ULIRGs trace the warm dense gas that is tightly correlated with the massive star forming activity. While the values of $T_\mathrm{warm}$ and \nhto\ (by assuming that they have similar velocity dispersion $\Delta V$) are similar to the local ones, $\tau_{100}$ in the \hz\ Hy/ULIRGs is likely to be greater than 1 (optically thick), which is larger than $\tau_{100}=0.05\text{--}0.2$ found in the local \ir\ galaxies. However, we notice that the parameter space is still not well constrained in our sources through \hto\ excitation modelling. Due to the limited excitation levels of the detected \hto\ lines, we are only able to perform the modelling with pure \fir\ pumping. The detection of relatively strong \htop\ lines opens the possibility to help understanding the formation of such large amount of \hto. In these \hz\ Hy/ULIRGs, the \hto\ formation is likely to be dominated by ion-neutral reactions powered by cosmic-ray-dominated regions. % The velocity-integrated flux density ratio between \htop\ and \hto\ ($I_{\mathrm{H_2O^+}}/\ihto \sim 0.3$), is remarkably constant from low to high-redshift, reflecting similar conditions in Hy/ULIRGs. However, more observations of \htop\ emission/absorption and also OH$^{+}$ lines are needed to further constrain the physical parameters of the cosmic-ray-dominated regions and the ionization rate in those regions. We have demonstrated that the submm \hto\ emission lines are strong and easily detectable with NOEMA. Being a unique diagnostic, the \hto\ emission offers us a new approach to constrain the physical conditions in the intense and heavily obscured star-forming regions dominated by \fir\ radiation at \hz. Follow-up observations of other gas tracers, for instance, CO, HCN, \htop\ and OH$^+$ using the NOEMA, IRAM 30m and JVLA will complement the \hto\ diagnostic of the structure of different components, dominant physical processes, star formation and chemistry in \hz\ Hy/ULIRGs. With unprecedented spatial resolution and sensitivity, the image from the ALMA long baseline campaign observation of SDP\,81 \citep[also known as {\it H}-ATLAS\,J090311.6+003906, ][]{2015ApJ...808L...4A, 2015MNRAS.452.2258D, 2015MNRAS.451L..40R}, shows the resolved structure of the dust, CO and \hto\ emission in the $z=3$ ULIRG. With careful reconstruction of the source plane images, ALMA will help to resolve the submm \hto\ emission in \hz\ galaxies into the scale of giant molecular clouds, and provide a fresh view of detailed physics and chemistry in the early Universe. \begin{acknowledgement} We thank our referee for the very detail comments and suggestions which have improved the paper. This work was based on observations carried out with the IRAM Interferometer NOEMA, supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain). The authors are grateful to the IRAM staff for their support. CY thanks Claudia Marka and Nicolas Billot for their help of the IRAM 30m/EMIR observation. CY also thanks Zhi-Yu Zhang and Iv\'an Oteo for insightful discussions. CY, AO and YG acknowledge support by NSFC grants \#11311130491, \#11420101002 and CAS Pilot B program \#XDB09000000. CY and YG also acknowledge support by NSFC grants \#11173059. CY, AO, AB and YG acknowledge support from the Sino-French LIA-Origin joint exchange program. E.G-A is a Research Associate at the Harvard-Smithsonian Center for Astrophysics, and thanks the Spanish Ministerio de Econom\'{\i}a y Competitividad for support under projects FIS2012-39162-C06-01 and ESP2015-65597-C4-1-R, and NASA grant ADAP NNX15AE56G. RJI acknowledges support from ERC in the form of the Advanced Investigator Programme, 321302, COSMICISM. US participants in {\it H}-ATLAS acknowledge support from NASA through a contract from JPL. Italian participants in {\it H}-ATLAS acknowledge a financial contribution from the agreement ASI-INAF I/009/10/0. SPIRE has been developed by a consortium of institutes led by Cardiff Univ. (UK) and including: Univ. Lethbridge (Canada); NAOC (China); CEA, LAM (France); IFSI, Univ. Padua (Italy); IAC (Spain); Stockholm Observatory (Sweden); Imperial College London, RAL, UCL-MSSL, UKATC, Univ. Sussex (UK); and Caltech, JPL, NHSC, Univ. Colorado (USA). This development has been supported by national funding agencies: CSA (Canada); NAOC (China); CEA, CNES, CNRS (France); ASI (Italy); MCINN (Spain); SNSB (Sweden); STFC, UKSA (UK); and NASA (USA). CY is supported by the China Scholarship Council grant (CSC No.201404910443). \end{acknowledgement} \footnotesize
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7
1607.06220
We report rest-frame submillimeter H<SUB>2</SUB>O emission line observations of 11 ultra- or hyper-luminous infrared galaxies (ULIRGs or HyLIRGs) at z 2-4 selected among the brightest lensed galaxies discovered in the Herschel-Astrophysical Terahertz Large Area Survey (H-ATLAS). Using the IRAM NOrthern Extended Millimeter Array (NOEMA), we have detected 14 new H<SUB>2</SUB>O emission lines. These include five 3<SUB>21</SUB>-3<SUB>12</SUB>ortho-H<SUB>2</SUB>O lines (E<SUB>up</SUB>/k = 305 K) and nine J = 2 para-H<SUB>2</SUB>O lines, either 2<SUB>02</SUB>-1<SUB>11</SUB>(E<SUB>up</SUB>/k = 101 K) or 2<SUB>11</SUB>-2<SUB>02</SUB>(E<SUB>up</SUB>/k = 137 K). The apparent luminosities of the H<SUB>2</SUB>O emission lines are μL<SUB>H<SUB>2</SUB>O</SUB> 6-21 × 10<SUP>8</SUP> L<SUB>⊙</SUB> (3 &lt;μ&lt; 15, where μ is the lens magnification factor), with velocity-integrated line fluxes ranging from 4-15 Jy km s<SUP>-1</SUP>. We have also observed CO emission lines using EMIR on the IRAM 30 m telescope in seven sources (most of those have not yet had their CO emission lines observed). The velocity widths for CO and H<SUB>2</SUB>O lines are found to be similar, generally within 1σ errors in the same source. With almost comparable integrated flux densities to those of the high-J CO line (ratios range from 0.4 to 1.1), H<SUB>2</SUB>O is found to be among the strongest molecular emitters in high-redshift Hy/ULIRGs. We also confirm our previously found correlation between luminosity of H<SUB>2</SUB>O (L<SUB>H<SUB>2</SUB>O</SUB>) and infrared (L<SUB>IR</SUB>) that L<SUB>H<SUB>2</SUB>O</SUB> L<SUB>IR</SUB><SUP>1.1-1.2</SUP>, with ournew detections. This correlation could be explained by a dominant role of far-infrared pumping in the H<SUB>2</SUB>O excitation. Modelling reveals that the far-infrared radiation fields have warm dust temperature T<SUB>warm</SUB> 45-75 K, H<SUB>2</SUB>O column density per unit velocity interval N<SUB>H<SUB>2</SUB>O</SUB> /ΔV ≳ 0.3 × 10<SUP>15</SUP> cm<SUP>-2</SUP> km<SUP>-1</SUP> s and 100 μm continuum opacity τ<SUB>100</SUB>&gt; 1 (optically thick), indicating that H<SUB>2</SUB>O is likely to trace highly obscured warm dense gas. However, further observations of J ≥ 4 H<SUB>2</SUB>O lines are needed to better constrain the continuum optical depth and other physical conditions of the molecular gas and dust. We have also detected H<SUB>2</SUB>O<SUP>+</SUP> emission in three sources. A tight correlation between L<SUB>H<SUB>2</SUB>O</SUB> and L<SUB>H<SUB>2</SUB>O<SUP>+</SUP></SUB> has been found in galaxies from low to high redshift. The velocity-integrated flux density ratio between H<SUB>2</SUB>O<SUP>+</SUP> and H<SUB>2</SUB>O suggests that cosmic rays generated by strong star formation are possibly driving the H<SUB>2</SUB>O<SUP>+</SUP> formation. <P />Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.The reduced spectra as FITS files 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/A80">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/595/A80</A>
false
[ "CO emission lines", "O", "H<SUB>2</SUB", "H<SUB>2</SUB>O column density", "CO and H<SUB>2</SUB>O lines", "their CO emission lines", "high redshift", "H-ATLAS", "ULIRGs", "dust", "anonymous ftp", "galaxies", "H<SUB>2</SUB>O", "strong star formation", "the H<SUB>2</SUB", "velocity-integrated line fluxes", "continuum opacity", "CO", "rest-frame submillimeter H<SUB>2</SUB>O emission line observations", "L<SUB>H<SUB>2</SUB>O</SUB" ]
11.657417
11.52552
-1
12409176
[ "Macías, Enrique", "Anglada, Guillem", "Osorio, Mayra", "Calvet, Nuria", "Torrelles, José M.", "Gómez, José F.", "Espaillat, Catherine", "Lizano, Susana", "Rodríguez, Luis F.", "Carrasco-González, Carlos", "Zapata, Luis" ]
2016ApJ...829....1M
[ "Imaging the Photoevaporating Disk and Radio Jet of GM Aur" ]
38
[ "Instituto de Astrofísica de Andalucía (CSIC), Glorieta de la Astronomía s/n, E-18008 Granada, Spain", "Instituto de Astrofísica de Andalucía (CSIC), Glorieta de la Astronomía s/n, E-18008 Granada, Spain", "Instituto de Astrofísica de Andalucía (CSIC), Glorieta de la Astronomía s/n, E-18008 Granada, Spain", "Department of Astronomy, University of Michigan, 825 Dennison Building, 500 Church Street, Ann Arbor, MI 48109, USA", "Institut de Ciències de l'Espai (CSIC)-Institut de Ciències del Cosmos (UB)/IEEC, Martí i Franquès 1, E-08028 Barcelona, Spain ;", "Instituto de Astrofísica de Andalucía (CSIC), Glorieta de la Astronomía s/n, E-18008 Granada, Spain", "Department of Astronomy, Boston University, 725 Commonwealth Avenue, Boston, MA 02215, USA", "Instituto de Radioastronomía y Astrofísica UNAM, Apartado Postal 3-72 (Xangari), 58089 Morelia, Michoacán, Mexico", "Instituto de Radioastronomía y Astrofísica UNAM, Apartado Postal 3-72 (Xangari), 58089 Morelia, Michoacán, Mexico", "Instituto de Radioastronomía y Astrofísica UNAM, Apartado Postal 3-72 (Xangari), 58089 Morelia, Michoacán, Mexico", "Instituto de Radioastronomía y Astrofísica UNAM, Apartado Postal 3-72 (Xangari), 58089 Morelia, Michoacán, Mexico" ]
[ "2017ApJ...834..138Z", "2017ApJ...838...97M", "2017ApJ...850..115S", "2017IAUS..328..264N", "2018A&A...610A..13C", "2018A&ARv..26....3A", "2018ASPC..517...87B", "2018ASPC..517..113B", "2018ASPC..517..137A", "2018ASPC..517..155P", "2018ASPC..517..281M", "2018ASPC..517..309G", "2018ASPC..517..345T", "2018ASPC..517..357H", "2018Ap&SS.363..246V", "2018ApJ...865...37M", "2018ApJ...865...77R", "2018arXiv181006598A", "2019A&A...624A...3P", "2019A&A...631A..58C", "2019ApJ...877L..34E", "2019ApJ...878...16P", "2019ApJ...884..159B", "2019BAAS...51c.129L", "2020A&A...642A.171R", "2020ApJ...891...48H", "2020ApJ...897...54W", "2020arXiv200702347V", "2021A&A...648A..33M", "2021A&A...650A.199M", "2021ApJ...913..122R", "2021ScPP...10...67S", "2022ApJ...924..104E", "2023A&A...673A..77R", "2023A&A...677A.118C", "2023ASPC..534..567P", "2023EPJP..138..225V", "2024A&A...684A.134R" ]
[ "astronomy" ]
19
[ "ISM: jets and outflows", "protoplanetary disks", "radio continuum: stars", "stars: individual: GM Aur", "stars: pre-main sequence", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1989ApJ...340L..69W", "1991ApJ...381..250B", "1995ApJ...452..736H", "1995ApJS..101..117K", "1995RMxAC...1...67A", "1998A&A...338L..63D", "2000ApJ...529..477L", "2001MNRAS.328..485C", "2002ApJ...581..357K", "2004ApJ...607..890F", "2004ApJ...614..807L", "2005ApJ...630L.185C", "2005MNRAS.358..283A", "2006A&A...446..211R", "2007IAUS..243..203C", "2008AJ....136.1852R", "2008MNRAS.391L..64A", "2009ApJ...690.1539G", "2009ApJ...697..957N", "2009ApJ...698..131H", "2009ApJ...702..724P", "2009ApJ...703.1203H", "2009ApJ...705.1237G", "2010A&A...519A.113G", "2010ApJ...717..441E", "2010MNRAS.401.1415O", "2010MNRAS.406.1553E", "2011A&A...532A..71R", "2011ARA&A..49...67W", "2011MNRAS.412...13O", "2012ApJ...747..142S", "2012ApJ...751...63A", "2012ApJ...751L..42P", "2012ApJ...752L..29C", "2013ApJ...762...62E", "2013ApJ...769...21S", "2013ApJ...771..129A", "2013MNRAS.434.3378O", "2014A&A...570L...9G", "2014ApJ...793L..21R", "2014ApJ...795....1P", "2014prpl.conf..475A", "2015ApJ...805..149I" ]
[ "10.3847/0004-637X/829/1/1", "10.48550/arXiv.1607.04225" ]
1607
1607.04225_arXiv.txt
\label{sec:intro} Photoevaporation, together with viscous accretion, is expected to play an important role in the dispersal of protoplanetary disks \citep{wil11,ale14}. High energy radiation -- i.e. far-UV (FUV), extreme-UV (EUV), and X-ray radiation -- originating at the stellar chromosphere of low-mass stars can ionize and heat the disk surface \citep{cla01,gor09,owe10}. Beyond a critical radius, the thermal energy of the heated surface becomes higher than its binding gravitational energy and the gas escapes in the form of a wind. While EUV photons produce a fully-ionized wind, X-rays can penetrate into deeper, neutral regions of the disk, creating a denser, partially-ionized photoevaporative wind \citep{gor09,owe11}. Although the early stages of disk evolution are dominated by viscous accretion, as the accretion rate decreases, central star-driven photoevaporation should eventually dominate over disk accretion, clearing the gas and leading the disk into the debris disk phase \citep{ale14}. The timescale of gas removal and, thus, the impact of photoevaporation in the disk evolution, will strongly depend on the ionization rate reaching the disk and the mass loss rate produced by the photoevaporative winds. So far, the primary diagnostic of disk photoevaporation has been optical and mid-IR forbidden line emission (e.g. [O I] $6300$ \AA ~and [Ne II] $12.81~\mu$m) from the wind \citep{fon04}. The redshifted side of the flow is blocked by the disk midplane, which is optically thick at these wavelengths. Therefore, the line profile is expected to be essentially narrow ($\sim10$ km s$^{-1}$) and blueshifted by 5--7 km s$^{-1}$ \citep{fon04}, although high disk inclinations and optically thin regions in the disk (like gaps or cavities) can produce broader lines centered at the systemic velocity \citep{ale08}. Blueshifted lines have been detected in a number of protoplanetary disks\citep{naj09,pas09,sac12}. However, similar line profiles and luminosities can be obtained with different models \citep{ale14}. Therefore, these lines cannot be used to constrain the high energy radiation responsible for the photoevaporative wind or to infer the mass loss rate in the flow \citep{erc10}. \citet{lug04}, and more recently \citet{ava12}, proposed that free-free emission at cm wavelengths could be used as a diagnostic of disk photoevaporation in massive stars. \citet{pas12} and \citet{owe13} followed a similar approach focusing on central star-driven photoevaporation in low-mass stars, and proposed that cm observations could actually provide a better observational test than forbidden line observations. The free-free emission from the fully (EUV case) or partially (X-rays case) ionized disk surface is optically thin and, thus, directly proportional to the ionizing radiation reaching the disk. Since the X-ray luminosity of T Tauri stars can be directly measured, one can in principle estimate the free-free emission produced by the X-ray-ionized gas and, therefore, estimate the EUV photon luminosity impinging on the disk from the remaining observed emission. Following this idea, recent observational studies at cm wavelengths have focused on the free-free emission of protoplanetary disks in order to constrain photoevaporation models \citep{gal14,pas14}. Due to their limited angular resolution, however, it is difficult to ascertain that these radio observations were not contaminated by free-free emission from an accretion-driven collimated jet. Since classical T Tauri stars present lower accretion rates than younger stellar objects, weak (or even absent) radio jet emission is expected in this type of sources. Nevertheless, thanks to the improved sensitivity of the Karl G. Jansky Very Large Array (VLA), \citet{rod14} were recently able to resolve the emission at 3.3 cm of a relatively weak radio jet in AB Aur, a Herbig Ae star surrounded by a transitional disk. GM Auriga is a well-known T Tauri star ($d\simeq$140 pc, K5 spectral type, $L_{\star}\simeq0.9$ $L_{\odot}$, $M_{\star}\simeq1.1~M_{\odot}$; \citealp{ken95}) surrounded by a transitional disk with a dust cavity of radius $\sim24$ au ($\sim0.17''$; \citealp{cal05,hug09,esp10}). [O I] and [Ne II] lines have been detected towards GM Aur, indicating the presence of high energy radiation reaching the disk. However, the [OI] spectrum has a very poor spectral resolution and the [NeII] spectrum shows no clear evidence of a blueshifted line peak that would confirm the presence of photoevaporation in GM Aur \citep{har95,naj09}. This could indicate that the lines are tracing a bounded ionized layer of the disk \citep{naj09}. Alternatively, the lack of an observed blueshifted line peak could actually be due to the disk cavity, which could allow the redshifted component of the wind to be visible \citep{owe13}, or due to an insufficient signal-to-noise ratio in the [Ne II] spectrum. Therefore, even though observations indicate that high energy radiation is impinging on the disk surface, the presence of photoevaporation in GM Aur is still uncertain. Here we report new sensitive high angular resolution VLA observations at 7 mm, 3 cm, and 5 cm towards the transitional disk of GM Aur, showing evidence of the presence of free-free emission from both photoevaporative winds and a radio jet.
\label{sec:conclusions} We have analyzed the results of multi-configuration VLA observations at Q, Ka, K, X, and C bands towards the transitional disk of GM Aur, revealing the presence of dust thermal and free-free emission at cm wavelengths. At 3 cm the emission presents an angularly resolved tripolar morphology that we separate into three components: the dust emission from the GM Aur disk, the free-free emission from a radio jet perpendicular to it, and the free-free emission from a photoevaporative wind arising from the disk. This is the first time that free-free emission from disk photoevaporation in a low mass star has been spatially resolved and separated from other components. We conclude that extreme-UV (EUV) radiation is the main agent responsible for the ionization of the photoevaporative wind in GM Aur, although requiring a low photon luminosity of $\Phi_{\rm EUV}\simeq6\times10^{40}$ s$^{-1}$. This low EUV photon luminosity produces a mass loss rate of only $\dot{M}_{w,EUV}\simeq1.3\times10^{-10}~M_{\odot}$ yr$^{-1}$. Therefore, other mechanisms, such as X-ray photoevaporation, are required to disperse the disk in the timescale imposed by observations. On the other hand, we estimate a mass loss rate in the radio jet in GM Aur of $\dot{M}_{{\rm out}}\simeq(3$-$5)\times10^{-9}~M_{\odot}$ yr$^{-1}$, which represents one of the lowest mass ejection rates in a jet estimated so far. Nevertheless, the ratio $\dot{M}_{out}/\dot{M}_{acc}\simeq0.1$, typical of younger protostars, seems to be valid as well for GM Aur. Therefore, our results suggest that disks with very low mass accretion rates still present collimated ejections of material, apparently following the same physical mechanisms as much younger protostars At least in GM Aur, the cm free-free emission of the jet and the photoevaporative wind seem to be of the same order. Future radio observations aiming to study photoevaporation in the last stages of disk evolution should be cautious and try to disentangle the contribution to the observed radio emission of the dust, the jet, and the photoevaporative winds.
16
7
1607.04225
Photoevaporation is probably the main agent for gas dispersal during the last stages of protoplanetary disk evolution. However, the overall mass-loss rate in the photoevaporative wind and its driving mechanism are still not well understood. Here we report multi-configuration Very Large Array observations at 0.7, 3, and 5 cm toward the transitional disk of GM Aur. Our radio continuum observations allow us to image and spatially resolve, for the first time, the three main components at work in this stage of the disk evolution: the disk of dust, the ionized radio jet perpendicular to it, and the photoevaporative wind arising from the disk. The mass-loss rate inferred from the flux density of the radio jet is consistent with the ratio between ejection and accretion rates found in younger objects, suggesting that transitional disks can power collimated ejections of material apparently following the same physical mechanisms as much younger protostars. Our results indicate that extreme-UV (EUV) radiation is the main ionizing mechanism of the photoevaporative wind traced by the free-free emission. The required low EUV photon luminosity of ∼6 × 10<SUP>40</SUP> s<SUP>-1</SUP> would produce a photoevaporation rate of only {\dot{M}}<SUB>w,{EUV</SUB>}≃ 1.3× {10}<SUP>-10</SUP> {M}<SUB>⊙ </SUB> yr<SUP>-1</SUP>. Therefore, other mechanisms are required to disperse the disk in the timescale imposed by observations.
false
[ "transitional disks", "protoplanetary disk evolution", "accretion rates", "GM Aur", "observations", "younger objects", "other mechanisms", "the transitional disk", "the disk evolution", "ejection", "the ionized radio jet perpendicular", "M}<SUB>⊙", "dust", "the disk", "the same physical mechanisms", "the main ionizing mechanism", "multi-configuration Very Large Array observations", "material", "gas dispersal", "first" ]
9.808196
12.933342
-1
2415187
[ "Querejeta, M.", "Schinnerer, E.", "García-Burillo, S.", "Bigiel, F.", "Blanc, G. A.", "Colombo, D.", "Hughes, A.", "Kreckel, K.", "Leroy, A. K.", "Meidt, S. E.", "Meier, D. S.", "Pety, J.", "Sliwa, K." ]
2016A&A...593A.118Q
[ "AGN feedback in the nucleus of M 51" ]
46
[ "Max Planck Institute for Astronomy, Königstuhl, 17, 69117, Heidelberg, Germany", "Max Planck Institute for Astronomy, Königstuhl, 17, 69117, Heidelberg, Germany", "Observatorio Astronómico Nacional, Alfonso XII, 3, 28014, Madrid, Spain", "Institut für theoretische Astrophysik, Zentrum für Astronomie der Universität Heidelberg, Albert-Ueberle-Str. 2, 69120, Heidelberg, Germany", "Departamento de Astronomía, Universidad de Chile, Camino del Observatorio 1515, Las Condes, Santiago, Chile; Centro de Astrofísica y Tecnologías Afines (CATA), Camino del Observatorio 1515, Las Condes, Santiago, Chile; Visiting Astronomer, Observatories of the Carnegie Institution for Science, 813 Santa Barbara St, Pasadena, CA, 91101, USA", "Max Planck Institute for Radioastronomy, Auf dem Hügel 69, 53121, Bonn, Germany", "CNRS, IRAP, 9 Av. Colonel Roche, BP 44346, 31028, Toulouse, France; Université de Toulouse, UPS-OMP, IRAP, 31028, Toulouse, France", "Max Planck Institute for Astronomy, Königstuhl, 17, 69117, Heidelberg, Germany", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH, 43210, USA", "Max Planck Institute for Astronomy, Königstuhl, 17, 69117, Heidelberg, Germany", "Physics Department, New Mexico Institute of Mining and Technology, 801 Leroy Place, Socorro, NM, 87801, USA", "Institut de Radioastronomie Millimétrique, 300 Rue de la Piscine, 38406 Saint Martin d', Hères, France; Observatoire de Paris, 61 Avenue de l'Observatoire, 75014, Paris, France", "Max Planck Institute for Astronomy, Königstuhl, 17, 69117, Heidelberg, Germany" ]
[ "2017A&A...599A..53I", "2017ApJ...835..174S", "2017ApJ...845...50N", "2017ApJ...846...71L", "2017ApJ...851...76M", "2017ApJ...851...98T", "2017MNRAS.470L.107J", "2018ASPC..517..457N", "2018ApJ...859...23N", "2018ApJ...859..144A", "2018ApJ...867..110B", "2018arXiv181007526N", "2019A&A...625A..19Q", "2019A&A...632A..33A", "2019A&A...632A..61G", "2019ApJ...875L...8R", "2019ApJ...877...74M", "2019MNRAS.482..194B", "2019MNRAS.483.4586F", "2019MNRAS.488..685M", "2019arXiv190407621M", "2020A&A...633A.134L", "2020AJ....160..167J", "2020ApJ...894...22Y", "2020MNRAS.492.3194S", "2020MNRAS.498.3633Z", "2020MNRAS.499....1F", "2021A&A...654A..90L", "2021IAUS..359..238M", "2021MNRAS.505.5469S", "2022A&A...659A.173E", "2022A&A...662A..89D", "2022EPJWC.26500010S", "2022MNRAS.515.1158G", "2022MNRAS.515.1705Z", "2023A&A...679A.115P", "2023A&A...680L..20S", "2023ApJ...943..168M", "2023ApJ...944L..15S", "2023RAA....23d5008V", "2023pcsf.conf...57E", "2024A&A...685A.122U", "2024MNRAS.528.1276M", "2024RAA....24f5006L", "2024arXiv240319843S", "2024arXiv240611398G" ]
[ "astronomy" ]
13
[ "galaxies: ISM", "galaxies: active", "galaxies: Seyfert", "galaxies: structure", "galaxies: jets", "Astrophysics - Astrophysics of Galaxies" ]
[ "1988ApJ...329...38C", "1992AJ....103.1146C", "1992pagn.conf..307W", "1993ARA&A..31..473A", "1994ApJ...421..122T", "1997AJ....113..225G", "1997ApJS..112..315H", "1998ApJ...493L..63S", "2000ApJ...529...93K", "2001ApJ...560..139T", "2002ApJ...567...97L", "2002ApJ...577...31C", "2002ApJ...579..530W", "2004ApJ...603..463B", "2004ApJ...613..898T", "2004ApJ...616L..55M", "2004ApJS..152...63G", "2005AAS...206.1401M", "2005sf2a.conf..721P", "2006MNRAS.365...11C", "2006SPIE.6270E..1VF", "2007A&A...461..143S", "2007A&A...467.1037K", "2007A&A...468..731J", "2007A&A...468L..49M", "2007AJ....133.1176H", "2007ApJ...664..204S", "2007ApJ...671..333K", "2008ApJ...681.1183K", "2009A&A...493..525I", "2009A&A...502..515G", "2009A&A...506..689I", "2009ApJ...698..198G", "2009ApJ...704..842B", "2010A&A...518L..42V", "2010A&A...518L.155F", "2010ApJ...710..903M", "2010MNRAS.402.2462C", "2010MNRAS.406..705B", "2011AJ....141...41D", "2011ApJ...728...29W", "2011ApJ...735...88A", "2011ApJ...740...42L", "2011MNRAS.418..591B", "2012A&A...539A...8G", "2012A&A...543A..99C", "2012ARA&A..50..455F", "2012ApJ...745L..34Z", "2012ApJ...757..136W", "2012MNRAS.425L..66M", "2013A&A...549A..51F", "2013A&A...558A.124C", "2013AJ....146...19L", "2013ARA&A..51..511K", "2013ApJ...762L..16M", "2013ApJ...764..117B", "2013ApJ...772..112S", "2013ApJ...774...12F", "2013ApJ...776...50C", "2013ApJ...777....5S", "2013ApJ...779...42S", "2013ApJ...779...43P", "2013ApJ...779...45M", "2013ApJ...779..173N", "2014A&A...562A..21C", "2014A&A...565A..46D", "2014A&A...567A.125G", "2014ApJ...780..186A", "2014ApJ...784....3C", "2014ApJ...784....4C", "2014ApJ...795...30B", "2014MNRAS.439..400Z", "2015A&A...574A..32G", "2015A&A...580A...1M", "2015A&A...580A..35G", "2015ARA&A..53...51S", "2015ApJ...799...26M", "2015ApJ...815..124H", "2015ApJS..219....5Q", "2015MNRAS.451.2517S", "2015MNRAS.452...32R", "2015PKAS...30..439M", "2015RAA....15..802F", "2016A&A...588A..33Q", "2016A&A...588A..41C", "2016ApJ...822L..26B", "2016MNRAS.455L..82L", "2016MNRAS.456.2861O" ]
[ "10.1051/0004-6361/201628674", "10.48550/arXiv.1607.00010" ]
1607
1607.00010_arXiv.txt
\label{Sec:introduction} \subsection{AGN feedback} \label{Sec:agnfeedback} Feedback from star formation (SF) and active galactic nuclei (AGN) plays a key role in reconciling cosmological simulations of galaxy formation and evolution with observations across different redshifts \citep{2015ARA&A..53...51S}. It is often invoked to explain the co-evolution of black holes and their host galaxies \citep{2013ARA&A..51..511K}, the mass-metallicity relation \citep{2004ApJ...613..898T,2008ApJ...681.1183K}, the bimodality in the colours of galaxies \citep{2015MNRAS.451.2517S}, the enrichment of the intergalactic medium (Martin et al. 2010), and it can prevent galaxies from over-growing in stellar mass \rev{relative to their dark matter halo} \citep[e.g.][]{2006MNRAS.365...11C}. By expelling molecular gas from the host galaxy, or by changing its ability to form stars, AGN feedback can also regulate star formation; it can either result in the suppression of star formation \citep{2014ApJ...780..186A}, or in its local enhancement \citep{2013ApJ...772..112S}. AGN feedback is necessary to alleviate the tension between simulations and observations for the most massive galaxies: for $M_* \gtrsim 10^{10}\,M_\odot$ cosmological simulations start to overpredict the stellar mass content of galaxies relative to the mass in their dark matter haloes \citep{2010ApJ...710..903M}. However, feedback is implemented in numerical models in a relatively \textit{ad hoc} way, adjusting its intensity so that the output stellar masses match observations; it remains to be confirmed to what extent these feedback levels are realistic. There is also some observational evidence for the relevance of AGN feedback in quenching star formation in massive $M_* > 10^{10}\,M_\odot$ galaxies \citep[e.g.][]{2016MNRAS.455L..82L}. \rev{The stellar mass of M51, the galaxy that we study here, is $\sim$7$\times 10^{10}\,M_\odot$ \citep{2015ApJS..219....5Q}, and it hosts a low-luminosity AGN. Feedback from the active nucleus might hold the key to the tight connection between bulge and black hole masses \citep[e.g.][]{2009ApJ...698..198G}; in other words, AGN feedback in M51 could be responsible for regulating bulge and black hole growth, preventing the galaxy from developing a massive, young bulge.} In a cosmological context, a number of studies have recently brought attention to the relevance of AGN-powered outflows \citep[see][for a review]{2012ARA&A..50..455F}; specifically, the last years have seen a plethora of detections of ionised ``winds'' driven by AGN activity, including valuable statistics based on large samples \citep[e.g.][]{2014ApJ...795...30B,2016arXiv160104715C}. There is also ample observational evidence of AGN-driven massive molecular outflows in relatively distant sources \citep[e.g.][]{2012A&A...543A..99C,2014A&A...562A..21C,2013A&A...549A..51F,2014A&A...565A..46D,2015A&A...580A..35G}; however, few nearby counterparts have been \rev{resolved}, and the details of feedback are still poorly understood. Notable exceptions AGN-driven include M51 \citep{2004ApJ...616L..55M,2007A&A...468L..49M,2015ApJ...799...26M}, NGC\,4258 \citep{2007A&A...467.1037K}, and since the advent of ALMA, IC\,5063 \citep{2015A&A...580A...1M} and NGC\,1068 \citep{2014A&A...567A.125G}. Careful analysis of the ALMA $0.5''$ (100\,pc) resolution CO(2-1) data of the radio galaxy IC\,5063 \citep{2015A&A...580A...1M} favours a scenario of a cold molecular gas outflow driven by an expanding radio plasma jet; this results in a high degree of \textit{lateral expansion}, in agreement with numerical simulations of radio jets expanding through a dense clumpy medium \citep[e.g.][]{2011ApJ...728...29W,2012ApJ...757..136W}. However, the high inclination of the disc of IC\,5063 ($i \sim 80^\circ$) complicates the analysis, and does not allow one to directly map the distribution of molecular gas relative to the jet. A massive (M$_{gas}\sim 3 \times 10^{7}$\,M$_{\odot}$) AGN-driven outflow of dense molecular gas has also been detected by ALMA using high density tracers in the inner 400\,pc of NGC\,1068, revealing \rev{complex kinematics} \citep{2014A&A...567A.125G}. These observations suggest that the outflow is efficiently regulating gas accretion in the circumnuclear disc ($r \lesssim 200$\,pc). \rev{Characterising the amount of dense (n(H$_2$)$>10^{4-5}$~cm$^{-3}$) molecular gas that is expelled through this process can help to construct a more complete multiphase picture of AGN feedback.} From a theoretical perspective, the mere existence of fast molecular outflows is problematic, as large velocities are expected to result in the dissociation of the molecular gas. Cooling into a two-phase medium has been proposed as a mechanism to explain how molecular gas survives outflows (Zubovas \& King 2012, 2014). The numerical simulations from Wagner \& Bicknell (2011) and \citet{2012ApJ...757..136W} also demonstrate the possibility of AGN feedback via radio plasma jets impinging on a clumpy ISM, showing how the interaction between jet and gas \rev{can result in significant lateral expansion (perpendicular to the direction of propagation of the radio jet)}. \subsection{The nucleus of M51} The grand-design spiral galaxy M51 constitutes a unique setup due to its proximity \citep[7.6~Mpc;][]{2002ApJ...577...31C} and the low inclination of the disc \citep[$i \sim 22^o$;][]{2014ApJ...784....4C}. Its well-studied Seyfert\,2 nucleus \citep{1997ApJS..112..315H,2011AJ....141...41D} is seen as two radio lobes that are filled with hot X-ray gas \citep{2001ApJ...560..139T} and an outflow of ionised gas \citep{2004ApJ...603..463B}. In the context of the AGN unification picture from \citet{1993ARA&A..31..473A}, the fact that only narrow lines are visible (and no broad lines), which determines the Seyfert\,2 nature of M51, would be explained by obscuration from a dusty torus almost perpendicular to our line of sight; the orientation of the radio jet \citep[inclined 15$^\circ$ with respect to the plane of the disc; ][]{1988ApJ...329...38C}, if perpendicular to the torus, as expected, supports this idea. We assume that the AGN location is given by the nuclear maser emission position determined by \citet{2015ApJ...815..124H}, RA=13:29:52.708, Dec=+47:11:42.810, which is less than $\sim$0.1$''$ away from the radio continuum peak \citep{1994ApJ...421..122T,2007AJ....133.1176H,2011AJ....141...41D}. Even if the potential impact of the AGN on the surrounding molecular material has been a matter of debate, high-resolution observations have demonstrated that both CO and HCN are participating in an outflow \citep{1998ApJ...493L..63S,2007A&A...468L..49M}, with an extraordinarily high HCN/CO ratio ($>2$) in the immediate vicinity of the AGN \citep{2015ApJ...799...26M}. In \citet{2015arXiv151003440Qalt}, we have studied the molecular gas flows across the full disc of M51, probing the transport of gas to the nucleus. Combining our stellar mass map \citep{2015ApJS..219....5Q} with the high-resolution CO gas distribution mapped by PAWS \citep{2013ApJ...779...42S, 2013ApJ...779...43P}, we have found evidence for gas inflow, with rates which are comparable to the amount of outflowing gas ($\sim$1\,$M_\odot$/yr), as we will show. Our goal is to understand the interplay between nuclear activity and the ISM in the nucleus of M51, relating it to the molecular gas inflow rates that we have already measured. Thus, we present a multi-wavelength study of the inner $\sim$1\,kpc of M51, the region affected by the radio plasma jet. In order to study the stratification in the response of the molecular gas to AGN feedback in M51, we have obtained new Plateau de Bure interferometric observations of dense gas tracers for the central $60''$ (2\,kpc) of the galaxy. We have detected and imaged three molecules in their $J=1-0$ transition (HCN, HCO$^+$, HNC), but here we will focus on the brightest one, HCN, and compare it to the bulk molecular gas traced by CO from PAWS. HCO$^+$ and HNC will be analysed in a forthcoming publication. The paper is structured as follows. We describe the new and archival data used in the analysis in Sect.\,\ref{Sec:data}. The main results are presented in Sect.\,\ref{Sec:results}, and discussed in Sect.\,\ref{Sec:discussion}. Finally, Sect.\,\ref{Sec:conclusions} consists of a summary, conclusions, and some open questions.
\label{Sec:conclusions} We have studied AGN feedback effects in a nearby spiral galaxy, M51, which hosts a low-luminosity active nucleus ($L_\mathrm{bol} \sim 10^{44}$\,erg\,s$^{-1}$) and a kpc-scale radio jet. The first important conclusion is that \textit{even with such a modest AGN}, the effects of feedback can be significant out to a distance of $\sim$500\,pc. The particular spatial configuration of M51 has allowed us to directly witness the interplay between the radio jet and molecular gas, because the galaxy is almost face-on and its radio jet is expanding through the disc, at least in the inner 1\,kpc (the jet has an inclination $\sim$15$^\circ$ with the plane of the galaxy). The area of the jet where optical ionised lines are detected, the ionisation cone, is largely depleted of molecular gas as traced by \mbox{CO(1-0)} at $1''$ resolution. Instead, CO emission seems to accumulate \textit{towards the edges} of the ionisation cone. This is an important result, as it indicates that molecular gas may not survive under the strong radiation field produced by the AGN, and questions the applicability of (bi)conical outflow models to more distant, unresolved molecular outflows. We find evidence for multiple components and disturbed kinematics in the molecular gas across the whole extent of the radio plasma jet. This becomes particularly clear when looking at the different kinematic response shown by CO and HCN, tracers of the bulk and dense phases of molecular gas, respectively. Therefore, relative differences between CO and HCN prove to be a useful diagnostic tool when it comes to probing feedback effects from radio jets. Mechanical shocks are the most likely explanation for the observed differences between both tracers. We also find increased turbulence (higher velocity dispersion) in the molecular gas across the whole region covered by the radio jet. The situation found in M51 is analogous to that recently observed in other nearby galaxies with similarly modest radio jets (e.g.~Krause~et~al.~2007; Morganti~et~al.~2015), and agrees with numerical simulations of radio jets expanding through a clumpy medium \citep{2011ApJ...728...29W,2012ApJ...757..136W}. Therefore, a new paradigm seems to be emerging, in which feedback through radio jets has complex implications for molecular gas (and, therefore, for star formation), probably pushing it in different directions and increasing its turbulence. It seems that outflows are not the only important consequences of AGN feedback; in addition to potential removal of molecular gas, injecting turbulence and therefore preventing molecular gas from forming new stars is an important part of the AGN response. Overall, we have shown that the feedback from the AGN in M51 is a multi-scale phenomenon. In addition to the large-scale impact on molecular gas across the radio jet area, the central $5''$ (180\,pc) display a more extreme version of feedback, which has been interpreted as a molecular outflow before \citep{2007A&A...468L..49M}. We have estimated the corresponding outflow rates with our data, $\dot{M}_{\mathrm{H}_2} \sim 0.9\,M_\odot/\mathrm{yr}$ and $\dot{M}_{\mathrm{dense}} \sim 0.6\,M_\odot/\mathrm{yr}$, and discussed geometrical caveats. It is worth noting that the typical velocities of this \rev{outflow are below the escape velocity for the whole galaxy, but they are sufficient to bring a substantial amount of molecular material out of the bulge. Therefore, in combination with the increased turbulence induced by the radio jet, the feedback from the AGN could be sufficient to prevent M51 from growing a massive, young bulge. This would explain the low star formation rates within the bulge, and it would contribute to maintaining the approximately constant relation between the mass of the bulge and that of the central supermassive black hole. We could be uncovering the mechanisms that permit black holes to regulate themselves in concert with the growth of their surroundings. } It would be important to confirm whether similar multi-scale mechanisms operate in more active galaxies, which do indeed have significant amounts of gas at velocities large enough to escape their host. However, this will prove to be observationally challenging, as those very active sources tend to be more distant, and therefore much longer integration times are required to achieve the same sensitivity with current-day interferometers; ALMA and NOEMA should be able to undoubtedly contribute in this direction. One of the important conclusions from our study is that both high spatial \textit{and spectral} resolution are necessary to obtain a complete picture of feedback effects, as multiple velocity components and kinematic differences between tracers can only be robustly characterised when sufficient velocity resolution is available. In a similar way, we have confirmed that the response of tracers of gas \rev{with different critical densities} (CO, HCN) are not redundant, and provide important hints as to what regions are impacted by the radio plasma jets. Therefore, future observations should go in the direction of high-resolution, multi-species observations of active nuclei, probing sufficiently large spatial regions to cover the various spatial effects involved. \small %
16
7
1607.00010
AGN feedback is invoked as one of the most relevant mechanisms that shape the evolution of galaxies. Our goal is to understand the interplay between AGN feedback and the interstellar medium in M 51, a nearby spiral galaxy with a modest AGN and a kpc-scale radio jet expanding through the disc of the galaxy. For this purpose, we combine molecular gas observations in the CO(1-0) and HCN(1-0) lines from the Plateau de Bure interferometer with archival radio, X-ray, and optical data. We show that there is a significant scarcity of CO emission in the ionisation cone, while molecular gas emission tends to accumulate towards the edges of the cone. The distribution and kinematics of CO and HCN line emission reveal AGN feedback effects out to r ~ 500 pc, covering the whole extent of the radio jet, with complex kinematics in the molecular gas which displays strong local variations. We propose that this is the result of the almost coplanar jet pushing on molecular gas in different directions as it expands; the effects are more pronounced in HCN than in CO emission, probably as the result of radiative shocks. Following previous interpretation of the redshifted molecular line in the central 5'' as caused by a molecular outflow, we estimate the outflow rates to be Ṁ<SUB>H<SUB>2</SUB></SUB> ~ 0.9 M<SUB>⊙</SUB>/ yr and Ṁ<SUB>dense</SUB> ~ 0.6 M<SUB>⊙</SUB>/ yr, which are comparable to the molecular inflow rates (~1 M<SUB>⊙</SUB>/ yr); gas inflow and AGN feedback could be mutually regulated processes. The agreement with findings in other nearby radio galaxies suggests that this is not an isolated case, and is probably the paradigm of AGN feedback through radio jets, at least for galaxies hosting low-luminosity active nuclei. <P />The reduced HCN(1-0) datacube is only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/593/A118">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/593/A118</A>
false
[ "molecular gas emission", "molecular gas", "other nearby radio galaxies", "molecular gas observations", "radio jets", "AGN feedback effects", "galaxies", "AGN feedback", "HCN line emission", "archival radio", "gas inflow", "CO emission", "optical data", "strong local variations", "AGN", "anonymous ftp", "the molecular gas", "the molecular inflow rates", "complex kinematics", "radiative shocks" ]
14.287346
7.662543
-1
12473391
[ "Wu, Meng-Ru", "Fernández, Rodrigo", "Martínez-Pinedo, Gabriel", "Metzger, Brian D." ]
2016MNRAS.463.2323W
[ "Production of the entire range of r-process nuclides by black hole accretion disc outflows from neutron star mergers" ]
153
[ "Institut für Kernphysik, Technische Universität Darmstadt, D-64289 Darmstadt, Germany", "Department of Physics, University of California, Berkeley, CA 94720, USA; Department of Astronomy and Theoretical Astrophysics Center, University of California, Berkeley, CA 94720, USA", "Institut für Kernphysik, Technische Universität Darmstadt, D-64289 Darmstadt, Germany; GSI Helmholtzzentrum für Schwerionenforschung, Planckstr. 1, D-64291 Darmstadt, Germany", "Department of Physics and Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA" ]
[ "2017ARNPS..67..253T", "2017ApJ...836..230C", "2017ApJ...836L..21N", "2017ApJ...846..170T", "2017ApJ...849..153F", "2017ApJ...850..179C", "2017ApJ...850L..37P", "2017ApJ...851L..45G", "2017CQGra..34j4001R", "2017CQGra..34o4001F", "2017JPhG...44h4007P", "2017LRR....20....3M", "2017MNRAS.468.1522V", "2017MNRAS.472..904L", "2017Natur.551...75S", "2017PASJ...69..102T", "2017PhRvD..95f3016C", "2017PhRvD..96d3001T", "2017PhRvD..96l3015W", "2017PhRvL.119w1102S", "2017Sci...358.1559K", "2017arXiv171005931M", "2018A&A...615A.132R", "2018ARNPS..68..237F", "2018ASSL..457..637G", "2018ApJ...852..109T", "2018ApJ...852L..29R", "2018ApJ...856..101M", "2018ApJ...856..138J", "2018ApJ...858...52S", "2018ApJ...860...89M", "2018ApJ...865...87O", "2018ApJ...866...51K", "2018ApJ...868...65W", "2018ApJ...869..130R", "2018ApJ...869L..35R", "2018EPJA...54..221L", "2018IJMPD..2742005H", "2018IJMPD..2743018F", "2018MNRAS.474.2419G", "2018MNRAS.478..692C", "2018MNRAS.478.3298W", "2018MNRAS.480.3871C", "2018MNRAS.481.3423W", "2018NatAs...2..980F", "2018PASA...35...17B", "2018PhRvD..97b3009K", "2018SSRv..214...62T", "2018arXiv180307873J", "2018pas7.conf..337W", "2019AIPC.2127b0013B", "2019AnPhy.41067923M", "2019AnPhy.41167958B", "2019ApJ...871..117S", "2019ApJ...874....5Z", "2019ApJ...875..106C", "2019ApJ...877L..24M", "2019ApJ...879...47E", "2019ApJ...880...81X", "2019ApJ...882...40J", "2019ApJ...882..163J", "2019ApJ...885...33H", "2019ApJ...885L..19C", "2019ApJ...886L..30N", "2019EPJA...55...50R", "2019EPJA...55..203S", "2019LRR....23....1M", "2019MNRAS.482.3373F", "2019MNRAS.486.2896S", "2019MNRAS.489.1697M", "2019MNRAS.489.5775C", "2019PhRvD.100b3008M", "2019PhRvD.100l4042E", "2019PhRvL.122f2701W", "2019PrPNP.107..109K", "2019supe.book..251T", "2020ApJ...889..171K", "2020ApJ...893..153K", "2020ApJ...894....9F", "2020ApJ...897...20Z", "2020ApJ...902L..34B", "2020JPhCS1668a2029N", "2020JPhCS1668a2044T", "2020MNRAS.491.1832P", "2020MNRAS.497.3221F", "2020PhR...886....1N", "2020PhRvC.102d5804G", "2020PhRvD.102j3001C", "2020PhRvD.102j3015G", "2020Physi...2..213T", "2021APPSB..31...18A", "2021ARep...65..385B", "2021Ap&SS.366..104D", "2021ApJ...906...94Z", "2021ApJ...906...98N", "2021ApJ...912..157R", "2021ApJ...913..100K", "2021ApJ...917...24Z", "2021ApJ...921..161D", "2021ApJ...922..269R", "2021JPlPh..87a8402A", "2021LRR....24....5K", "2021MNRAS.505.1661B", "2021MNRAS.505.5862W", "2021MNRAS.506.1850J", "2021MNRAS.506.3560G", "2021RvMP...93a5002C", "2021Sci...371..945C", "2021arXiv211113005C", "2021arXiv211215470A", "2022A&A...663A..70F", "2022A&A...665A..10L", "2022APPSB..32...19W", "2022ApJ...925...22P", "2022ApJ...925...43Q", "2022ApJ...933...22K", "2022ApJ...936...84R", "2022ApJ...941..156C", "2022ApJS..260...27R", "2022CQGra..39a5008N", "2022EPJA...58...99C", "2022FrASS...9.4980V", "2022MNRAS.510.2820J", "2022MNRAS.512.4948B", "2022MNRAS.515..631G", "2022MNRAS.516.4760C", "2022NatRP...4..306S", "2022PhRvD.105h3024J", "2022arXiv221207498J", "2023A&ARv..31....1A", "2023AJ....165..100J", "2023ApJ...944..123V", "2023ApJ...951L..12J", "2023ApJS..268...66R", "2023CQGra..40h5008H", "2023EPJA...59...12T", "2023EPJA...59...67M", "2023MNRAS.519.1349W", "2023MNRAS.520.2727N", "2023MNRAS.523.2551K", "2023MNRAS.524.5514N", "2023MNRAS.525.3384K", "2023MNRAS.525.6249O", "2023MNRAS.526..952F", "2023PhRvD.108l3038A", "2023arXiv230209188F", "2023arXiv230306366K", "2023arXiv230804485B", "2024ApJ...961..119F", "2024MNRAS.529.2918G", "2024MNRAS.530.2336R", "2024MNRAS.531.4422G", "2024arXiv240511234M" ]
[ "astronomy", "physics" ]
8
[ "accretion", "accretion discs", "dense matter", "gravitational waves", "neutrinos", "nuclear reactions", "nucleosynthesis", "abundances", "stars: neutron", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics", "General Relativity and Quantum Cosmology", "Nuclear Theory" ]
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[ "10.1093/mnras/stw2156", "10.48550/arXiv.1607.05290" ]
1607
1607.05290_arXiv.txt
\label{sec-intro} Approximately half of the elements with mass number $A > 70$, and all of the transuranic nuclei, are formed by the rapid neutron capture process (the $r$-process; \citealt{Burbidge+57}, \citealt{Cameron57}). The astrophysical site of this process has been under debate for more than 50 yrs (see, e.g., \citealt{Qian&Wasserburg07}, \citealt{Arnould+07}, \citealt{Sneden+08}, \citealt{Thielemann+11} for reviews). Neutrino-driven outflows from proto neutron stars (NSs) following core collapse supernovae have for long been considered the prime candidate site (\citealt{Meyer+92}; \citealt{Woosley+94}; \citealt{Qian&Woosley96}). However, state-of-the-art calculations find thermodynamic conditions that are at best marginal for the $r$-process, especially when extending up to the heaviest third-peak elements with mass number $A \sim 195$ (e.g.~\citealt{MartinezPinedo+12}; \citealt{Roberts+12}; \citealt{martinez2014}). Prospects for a successful $r$-process in neutrino-driven outflows may be improved if the proto-NS is born with a strong magnetic field and very rapid rotation (e.g.~\citealt{Thompson+04,Metzger+07,Vlasov+14}). If the supernova itself is MHD-powered, additional magnetocentrifugal acceleration could substantially reduce the electron fraction of the outflow compared to its value in the purely neutrino-driven case favouring the occurrence of an $r$-process (e.g.~\citealt{Burrows+07}, \citealt{Winteler.Kaeppeli.ea:2012}, \citealt{Nishimura+15}). However, current simulations of MHD-powered supernova explosions need further improvements, especially considering the role of instabilities on the jet structure which manifest in three dimensions (\citealt{Moesta+14}). The coalescence of double NS (NS--NS) and NS-black hole (NS--BH) binaries \citep{Lattimer&Schramm74} provides an alternative $r$-process source. Numerical simulations of these events show that a robust outcome of the merger is the ejection of $\sim 10^{-4}$--$10^{-1}$~$M_\odot$ of highly neutron-rich matter on the dynamical time (e.g.~\citealt{Hotokezaka+13}, \citealt{Bauswein+13}; see \citealt{lehner2014} and \citealt{Baiotti:2016qnr} for recent reviews). Estimates show that NS--NS/NS--BH mergers could contribute a sizable fraction of the total production of $r$-process elements in the Galaxy, depending on the uncertain merger rates. At the same time, previous arguments against mergers being dominant $r$-process sites based on Galactic chemical evolution and the observed prompt enrichment of $r$-process nuclei in metal poor stars (e.g.~\citealt{Argast+04}) have been challenged (e.g., \citealt{vdvoort2015,shen2015,hirai2015}). Additional evidence supporting the presence of a `high yield' $r$-process site -- like an NS--NS/NS--BH merger -- includes the discovery of highly $r$-process enriched stars in the ultra-faint dwarf galaxy Reticulum II (\citealt{Ji+16}), and the abundance of the short-lived isotope $^{244}$Pu on the sea floor (\citealt{Wallner+15, Hotokezaka+15}). Nucleosynthesis in NS--NS/NS--BH mergers has also received a recent surge of interest due to the realization that the radioactive decay of the $r$-process ejecta can power a thermal transient (a ``kilonova''; e.g.,~\citealt{Li&Paczynski98}, \citealt{Metzger+10}, \citealt{Roberts+11}, \citealt{Barnes&Kasen13}, \citealt{tanaka2013}), which could serve as a promising electromagnetic counterpart to the gravitational waves (\citealt{Metzger&Berger12}). The detection of a possible kilonova following the {\it Swift} GRB 130603B (\citealt{Berger+13}; \citealt{Tanvir+13}) highlights the potential of kilonovae as both a unique diagnostic of physical processes at work during the merger and a direct probe of the formation of $r$-process nuclei (see, e.g., \citealt{rosswog2015}, \citealt{FM16}, \citealt{tanaka2016} for recent reviews). Previous work on the $r$-process in NS--NS/NS--BH mergers has been focused primarily on the dynamical ejecta that is unbound promptly during the immediate aftermath of the merger (e.g.~\citealt{Meyer89}, \citealt{Freiburghaus+99}, \citealt{goriely2005}). Earlier simulations that did not include weak interactions have shown this unbound matter to be highly neutron-rich, with an electron fraction $Y_e \lesssim 0.1$, sufficiently low to produce a robust abundance pattern for heavy nuclei with $A \gtrsim 130$ as the result of fission cycling (e.g., \citealt{Goriely+11}, \citealt{Korobkin+12}, \citealt{Bauswein+13}, \citealt{Mendoza-Temis.Wu.ea:2015}). More recently, a number of merger calculations that include the effects of $e^\pm$ captures and neutrino irradiation in full general-relativity have shown that the dynamical ejecta can have a wider electron fraction distribution($Y_e \sim 0.1-0.4$) than models without weak interaction effects~\citep{sekiguchi2015,foucart2015a,radice2016}. As a result, lighter $r$-process elements with $90 \lesssim A \lesssim 130$ are generated in addition to third-peak elements \citep{wanajo2014}. It is important to keep in mind, however, that the light element yields in these calculations are dependent on the assumed dense-matter equation of state and on the details of the neutrino transport employed, in addition to the NS radii and the binary mass ratio. In addition to ejecting material dynamically, NS--NS/NS--BH mergers result in the formation of an accretion disc around the central remnant (e.g., \citealt{Oechslin&Janka06}); with the latter being a promptly formed BH or a longer-lived hypermassive NS (HMNS) (e.g., \citealt{shibata2000}). In both cases, the accretion disc can generate outflows on time-scales much longer than the orbital time (e.g., \citealt{Metzger+08, Lee+09,Metzger+09a}), and with a contribution to the total mass ejection that can be comparable to, or even larger than that from the dynamical ejecta (Fig/~\ref{f:brian_ejecta_mass}, see also \citealt{FM16}). A relatively massive disc ($\sim 0.1M_{\odot}$) can be formed following a NS--NS merger, as part of the process by which the HMNS sheds angular momentum outwards prior to collapsing into a BH (e.g.~\citealt{shibata2006}). Long-term hydrodynamic simulations of the disc evolution show that a significant fraction of the initial disc mass ($\sim 5-20\%$, corresponding to $\sim 0.01M_{\odot}$) is unbound in outflows powered by heating from angular momentum transport and nuclear recombination, on a timescale of $\gtrsim 1$~s (\citealt{Fernandez&Metzger13}, hereafter FM13; \citealt{just2014}, \citealt{FKMQ14}). As the result of weak interactions, the electron fraction of the disc outflows lies in the range $Y_e \sim 0.2-0.4$, generally higher than that of the dynamical ejecta, but still sufficiently low to achieve the $r$-process \citep{just2014}. \begin{figure} \includegraphics*[width=\columnwidth]{brian_ejecta_mass.pdf} \caption{Mass ejected dynamically during a compact binary merger versus that ejected in disc outflows. Each point corresponds to the result of a single time-dependent NS--NS (triangles) and BH-NS (squares) simulation. Shown are models by \citet{Hotokezaka+13} (blue), \citet{Oechslin&Janka06} (green, upper limits shown by arrows), \citet{just2014} (brown), \citet{East+12} (red), and \citet{foucart2014} (orange). The mass unbound in disc outflows is estimated to be 10 per cent of the mass of the remnant disc, based on calculations of the subsequent accretion disc evolution (e.g., FM13). Dashed lines show total ejecta mass contours (dynamical + disc winds) of $0.01 M_{\odot}$ and $0.1M_{\odot}$, bracketing the range necessary to explain the Galactic production rate of heavy $r$-process nuclei $\sim 5\times 10^{-7}M_{\odot}$~yr$^{-1}$ \citep{Qian00}, given the allowed range of the rates of NS--NS mergers $\in [4,61]$~Myr$^{-1}$ (99$\%$ confidence) calculated based on the population of Galactic binaries \citep{kim2015}. In reality, the ejecta mass range required to reproduce the Galactic abundances is uncertain by greater than an order of magnitude, due to systematic uncertainties on the merger rate and depending on the precise atomic mass range under consideration (e.g.,~\citealt{Bauswein:2014vfa}). See also \citet{FM16}.} \label{f:brian_ejecta_mass} \end{figure} Although most of the previous work on merger disc wind nucleosynthesis has focused on parametrized outflows powered by neutrino heating (e.g.~\citealt{McLaughlin&Surman05,Surman+08,Caballero+12,Surman+14}), in analogy with proto-NS winds, time-dependent models of the long-term disc evolution show that neutrino heating is sub-dominant relative to viscous heating in driving most of the disc outflow when a BH sits at the centre (FM13, \citealt{just2014}). Neutrino heating is much more important if the merger produces a long-lived HMNS \citep{Dessart+09,Metzger&Fernandez14,perego2014,martin2015}, resulting in a larger ejecta mass with higher electron fraction, depending on the uncertain lifetime of such a remnant. To date, the only fully time-dependent nucleosynthesis study of long-term outflows from BH accretion discs was carried out by \citet{just2014}, who found that disc outflows can generate elements from $A\sim 80$ to the actinides, with the contribution above $A =130$ being sensitive to system parameters. In this paper we further investigate nucleosynthesis in the outflows from NS--NS/NS--BH merger remnant accretion discs around BHs, by applying a nuclear reaction network on thermodynamic trajectories extracted from fully time-dependent, long-term hydrodynamic simulations of disc outflows. Our aim is to carry out a systematic study of the dependence of the $r$-process production on system parameters such as disc mass or viscosity, and on additional ingredients such as nuclear physics inputs to the reaction network or the feedback from nuclear heating on the disc dynamics. Our main conclusion is that disc outflows from NS binary mergers can in principle produce both the light and heavy $r$-process elements, without necessarily requiring additional contributions from the dynamical ejecta. The paper is organized as follows. Section \ref{sec-methods} describes the hydrodynamic models, extraction of thermodynamic trajectories, and the properties of the nuclear reaction network. Results and analysis are presented in section \ref{sec-results}. Finally, section \ref{sec-discussion} summarizes our findings and discusses broader astrophysical implications.
\label{sec-discussion} We have studied the production of $r$-process elements in the disc outflows from remnant accretion tori around BHs formed in NS-NS/NS-BH mergers. We used tracer particles in long-term, time dependent hydrodynamic simulations of these discs to record thermodynamic and kinematic quantities. The resulting trajectories were then post-processed with a dynamic $r$-process nuclear reaction network. Our results can be summarized as follows: \newline \noindent 1. -- Outflows from merger remnant discs around BHs can robustly generate light $r$-process elements with $A \lesssim 130$, regardless of the astrophysical parameters of the disc or the nuclear physics inputs employed (Fig.~\ref{fig-intd_sep_sdef}, Table~\ref{tab:models-nuc}). \newline \noindent 2. -- The yield of elements with $A > 130$ is most sensitive to the type and magnitude of the angular momentum transport process (Fig.~\ref{fig-intd_model_comp}d) and on the nuclear physics inputs employed in the nuclear reaction network (Fig.~\ref{fig-nucinput}). Given the physics employed in our hydrodynamic models, the $\alpha$ viscosity parameter is a key factor determining the abundance of third peak and heavier elements. Other parameters such as the disc mass or initial entropy have a relatively smaller impact on the abundances. \newline \noindent 3. -- We have identified a spike in the abundance of $A=132$ elements that arises whenever the disc outflow is highly convective, as is the case when using reasonable choices for the disc parameters (Fig.~\ref{fig-intd_sep_sdef}f). This feature can be erased if the disc evolution is fast or if the heating rate in the disc is very low, so that convection is suppressed. \newline \noindent 4. -- Inclusion of energy deposition from charged-particle reactions beyond $^4$He recombination can affect the ejecta dynamics and nucleosynthesis to the point where the $A=132$ abundance anomaly disappears (Fig.~\ref{fig-intd_s_h}). Alternatively, the processes responsible for controlling angular momentum transport and the thermodynamics of the disc (e.g. MHD and neutrino transport) can have a sensitive nucleosynthetic impact. The properties of convective motions in an hydrodynamical $\alpha-$disc may differ substantially from those of MHD turbulence (\citealt{Balbus&Hawley02}). \newline \noindent 5. -- The comparison with abundances observed in metal-poor stars shows that if the disc outflows contribute dominantly to the NS-NS/BS-BH ejecta, different initial configurations of the disc may account for the variation of light-to-heavy abundance ratio seen in these stars. \newline Our results together with those of~\citet{just2014} show that disc outflows are fundamental to understand $r$-process nucleosynthesis in mergers. For cases in which little dynamical ejecta is generated, {\it the disc outflow alone can contribute substantially to the heavy $r$-process element enrichment}, even while producing proportionally more elements with $A < 130$. This result reinforces the general view that NS-NS/NS-BH mergers are the primary astrophysical site for heavy $r$-process elements. It is also important because the presence of high-$Y_e$ dynamical ejecta in NS-NS mergers is uncertain theoretically, due in part to the dependence of shock-heated polar ejecta on the NS radius and equation of state. The early ejection of material with high $Y_e$ in the disc outflow has the potential to generate a blue peak in the kilonova, since it will generally reside on the outer layers of the ejecta \citep{FKMQ14,KFM15}. This nonetheless depends on the viewing directions close to the rotation axis to be relatively free of Lanthanide-rich dynamical ejecta material, which is usually the case for BH-NS mergers \citep{FQSKR15,KFM15}. The net kilonova contribution of systems studied is Lanthanide-rich, as inferred from Table~\ref{tab:models-nuc}. The work of \citet{kasen2013} shows that even a mass fraction of $\sim 10^{-2}$ in Lanthanides can increase the optical opacity by at least an order of magnitude relative to that of iron-like elements. Nearly all of our models have a Lanthanide mass fraction bigger than $1\%$ and thus while an outer Lanthanide-free ``skin" is usually obtained, the bulk of the wind will lead to an infrared transient. More promising in this respect is the possible onset of a neutron-powered precursor if a small fraction of material escapes quickly enough to freeze-out before neutrons are captured by heavy seeds \citep{metzger2015_n-precursor}. Simulations at much higher resolution than currently available are needed to resolve this question. A long-lived NS remnant can also increase the quantity of high-$Y_e$ ejecta, producing a more prominent blue component of the kilonova (\citealt{Metzger&Fernandez14}). A crucial improvement to our calculation would involve obtaining trajectories from an accretion disc outflow in which angular momentum transport -- and the associated energy dissipation -- is carried out by MHD stresses. Such a calculation can differ from ours in a number of ways. First, the change in the entropy due to viscous heating in an $\alpha$-viscosity model is likely different in an MHD disc, with the associated change in the equilibrium $Y_e$ that the weak interactions try to achieve. Secondly, the amount of mass ejected and the associated $Y_e$ distribution can change, altering the relative amounts of heavy- and light $r$-process elements in the outflow composition. Finally, the kinematic properties of the wind can change, in particular the velocity, which controls the expansion time, and the angular distribution, which is associated with the level of neutrino irradiation of the ejecta (material on the equatorial plane is more effectively shadowed from neutrinos than material leaving at high latitudes). Our calculations would also benefit from better neutrino transport, although the magnitude of the difference introduced can be comparable to that due to the spin of the BH (e.g., compare the results of \citealt{FKMQ14} and \citealt{just2014}). A more thorough, self-consistent treatment of nuclear heating would also make our calculations more realistic. In this respect, it is worth noting that the results of \citet{just2014} do not appear to show the abundance spike at $A=132$ that we obtain in many of our models. We surmise that this difference arises due to their inclusion of a single species of heavy nucleus ($^{54}$Mn) in the equation of state, which partially accounts for the energy production beyond $\alpha$ formation. Finally, it is important to consider the combined evolution of disc and dynamical ejecta in computing the net $r$-process yield from NS-NS/NS-BH mergers. Part of the dynamical ejecta is gravitationally bound, and mixes with the accretion disc, increasing the neutron-richness of the disc and therefore lowering the peak of the $Y_e$ distribution of the disc outflow \citep{FQSKR15}.
16
7
1607.05290
We consider r-process nucleosynthesis in outflows from black hole accretion discs formed in double neutron star and neutron star-black hole mergers. These outflows, powered by angular momentum transport processes and nuclear recombination, represent an important - and in some cases dominant - contribution to the total mass ejected by the merger. Here we calculate the nucleosynthesis yields from disc outflows using thermodynamic trajectories from hydrodynamic simulations, coupled to a nuclear reaction network. We find that outflows produce a robust abundance pattern around the second r-process peak (mass number A ∼ 130), independent of model parameters, with significant production of A &lt; 130 nuclei. This implies that dynamical ejecta with high electron fraction may not be required to explain the observed abundances of r-process elements in metal poor stars. Disc outflows reach the third peak (A ∼ 195) in most of our simulations, although the amounts produced depend sensitively on the disc viscosity, initial mass or entropy of the torus, and nuclear physics inputs. Some of our models produce an abundance spike at A = 132 that is absent in the Solar system r-process distribution. The spike arises from convection in the disc and depends on the treatment of nuclear heating in the simulations. We conclude that disc outflows provide an important - and perhaps dominant - contribution to the r-process yields of compact binary mergers, and hence must be included when assessing the contribution of these systems to the inventory of r-process elements in the Galaxy.
false
[ "black hole accretion discs", "compact binary mergers", "neutron star", "nuclear physics inputs", "angular momentum transport processes", "metal poor stars", "Disc outflows", "disc outflows", "mass number A", "nuclear heating", "nuclear recombination", "neutron star-black hole mergers", "hydrodynamic simulations", "double neutron star", "r-process elements", "initial mass", "r-process nucleosynthesis", "the Solar system r-process distribution", "model parameters", "A" ]
6.023593
2.360842
62
12436410
[ "Barrau, Aurélien", "Grain, Julien" ]
2016arXiv160707589B
[ "Cosmology without time: What to do with a possible signature change from quantum gravitational origin?" ]
21
[ "-", "-" ]
[ "2016arXiv161007467B", "2016arXiv161201236A", "2016arXiv161201964C", "2017CRPhy..18..207W", "2017PhRvD..95d6001C", "2017PhRvD..96b6002B", "2017PhRvD..96f6025B", "2017arXiv170503772B", "2018CQGra..35j5017H", "2018IJMPD..2750067M", "2018JCAP...05..072B", "2018PhRvD..98h6015S", "2020IJMPA..3530007O", "2020Univ....6...39A", "2020Univ....6..125B", "2021FrASS...8...81N", "2022JETPL.116...54B", "2022Univ....8..497B", "2022arXiv220407828B", "2023PhRvD.107l6008D", "2024PhRvD.109b4040F" ]
[ "astronomy", "physics" ]
6
[ "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]
[ "1959PhRv..116.1322A", "1994CQGra..11..389L", "2005CQGra..22.3349A", "2010IJMPD..19.1119S", "2010PhRvD..82h4035B", "2011CQGra..28u3001A", "2011arXiv1102.3660R", "2012CQGra..29h5005W", "2012CQGra..29h5009M", "2012CQGra..29i5010C", "2012CQGra..29u5013W", "2012PhRvD..85l3534C", "2012PhRvD..86h7301C", "2012PhRvL.109y1301A", "2013CQGra..30h5014A", "2013PhRvD..87d3507A", "2013PhRvD..87d4044B", "2013PhRvD..87j7503L", "2014CQGra..31e3001B", "2014EL....10840003M", "2014IJMPD..2342026R", "2014JCAP...03..048M", "2014PhLB..739..405B", "2014PhRvD..90l7503B", "2014arXiv1411.0272M", "2015FrP.....3...33B", "2015JCAP...05..051B", "2015JCAP...08..052B", "2015PhRvD..91h4035B", "2015PhRvD..92d5043B", "2015PhRvD..92f5002B", "2015PhRvD..92j4020H", "2015arXiv151103702F", "2016IJMPD..2542003G", "2016IJMPD..2542008B", "2016JCAP...02..022B", "2016PhRvD..93b3531S", "2016PhRvD..93l4011B", "2017PhLB..772...58B" ]
[ "10.48550/arXiv.1607.07589" ]
1607
1607.07589_arXiv.txt
\label{sec:intro} As a theory of the dynamics of space-time, General Relativity (GR) is most commonly phrased in a Lagrangian framework, where the gravitational fields is described by the metric tensor of (pseudo)riemanian manifolds. In this description, the "time" part of space-time is easily figured out from the signature of the metric tensor. Basically, Einstein's field equations are the equations of motion of a metric tensor which has the Lorentzian signature, $(-,+,+,+)$, and physical space-times are thus in the set of four dimensional, pseudoriemanian manifolds\footnote{The prefix "pseudo" is used to denote Lorentzian signature.}. The fact that a minus sign in the signature traces the presence of a "time" part in physical space-times can be intuitively understood in a causal sense. Because of that minus sign, the interval separating two space-time events can be either space-like (i.e. positive-valued), light-like (i.e. null), or time-like (i.e. negative-valued). Two events which are separated by a space-like interval are causally disconnected since no geodesic traced by usual matter and radiation fields can link them. Roughly said, the distance between these two space-time events is such that no physical fields travels fast enough to cover that distance within the time interval separating the two events. For light-like and time-like separations however, physical fields can join the two events. Clearly, if the signature is Euclidean, $(+,+,+,+)$, such an interpretation is not possible since intervals are all positive-valued. \\ In the Hamiltonian language of GR, the signature of space-time is traced through the so-called algebra of hypersurface deformation. The Hamiltonian formulation of GR, as firstly proposed by Arnowitt, Deser and Misner \cite{PhysRev.116.1322}, relies on the foliation of 4-dimensional (pseudo)riemanian manifolds into a set of 3-dimensional hypersurfaces, $\left(\Sigma_t\right)_{t\in\mathbb{R}}$ (see \cite{lqg6} for details). On a given hypersurface, the canonical variables are the induced metric, $q_{ab}$, and its canonically conjugate momentum, $P^{ab}:=\frac{\delta S}{\delta \dot{q}_{ab}}$, with $\dot{q}_{ab}$ the derivative with respect to $t$ of the induced metric. In such a setting, the dynamics of the gravitational fields is the evolution of this set of canonical variables from one hypersurface to another. This phase space has however to be extended to include the lapse function, $N$, the shift vector, $N^a$, and their associated momenta, dubbed $C=\frac{\delta S}{\delta \dot{N}}$ and $C_a=\frac{\delta S}{\delta \dot{N}^a}$ respectively. This describes the adopted foliation, and since GR is independent of such a choice through diffeomorphism invariance, their is no evolution of $N$ and $N^a$, i.e. the Einstein-Hilbert action does not depend on $\dot{N}$ and $\dot{N}^a$. As a result, the momenta associated to $N$ and $N^a$ are constrained to zero and this should be preserved through evolution. This in turn shows that $N$ and $N^a$ are mere Lagrange multipliers and the Hamiltonian reads $\mathcal{H}=\int\dd^3x\left[NH(q_{ab},P^{ab})+N^aH_a(q_{ab},P^{ab})\right]$ ; $H$ and $H_a$ are the Hamiltonian and spatial diffeomorphism constraints. Because this holds for any choice of the lapse function and the shift vector, $H\approx0$ and $H_a\approx0$ on solutions of GR, meaning that GR is a totally constrained system as a result of diffeomorphism invariance. This property of the theory can be phrased into an algebraic structure. One can define the smeared constrained, $S(N)=\int\dd^3xNH$ and $D(N^a)=\int\dd^3xN^aH_a$, and their time evolution is given by the Poisson brackets $\left\{\mathcal{H},S\right\}$ and $\left\{\mathcal{H},D\right\}$. Solutions of GR are such that $H=H_a=0$ and this should obviously be preserved across evolution. This can be shown to be equivalent to an algebraic structure satisfying the constrained themselves: \begin{align} &\left\{D(M^a),D(N^a)\right\}=D\left(M^b\partial_bN^a-N^b\partial_bM^a\right), \label{eq:algebra1} \\ &\left\{D(M^a),S(N)\right\}=S\left(M^b\partial_bN-N\partial_bM^b\right), \label{eq:algebra2} \\ &\left\{S(M),S(N)\right\}=sD\left(q^{ab}(M\partial_bN-N\partial_bM)\right),\label{eq:algebra3} \end{align} with $s=1$ for the Lorentzian signature, and $s=-1$ for the Euclidean one \cite{lqg6}. This means that the subspace of the phase space which satisfies the constraints, called the surface of constraints, is preserved through evolution. The constraints, also denoted $\mathcal{C}_I$ with $\mathcal{C}_1\equiv D$ and $\mathcal{C}_2\equiv S$, are said to form a system of first-class, i.e. $\left\{\mathcal{C}_I,\mathcal{C}_J\right\}={f^K}_{IJ}(q_{ab},P^{ab}) \mathcal{C}_K$ with ${f^K}_{IJ}$ structure functions depending on the phase-space variables. Interestingly enough, the signature explicitly appears in the last Poisson bracket of the algebra of constraints. This is in fact not a surprise since the above Hamiltonian description admits a clear geometrical interpretation \cite{HOJMAN197688}. Solution of GR are 4-dimensional (pseudo)riemanian manifolds. In the Hamiltonian framework, one embeds hypersurfaces in this 4-dimensional structure. The embedding defines two possible way of deforming the embedded hypersurfaces: one can either displace a point {\it within} the hypersurface, or one can displace a point in the direction orthogonal to the considered hypersurface. These two transformations are described by generators, $\mathcal{D}_a$ and $\mathcal{D}$, which respectively generates diffeomorphisms preserving $\Sigma_t$, and diffeomorphisms orthogonal to $\Sigma_t$. The two generators form an algebraic structure called the {\it algebra of hypersurface deformation} that is exactly the above algebra formed by the constraints themselves. In other words, the Hamiltonian and spatial diffeormorphism constraints are {\it representations} of the transformations of the embedded hypersurfaces. And, in fact, this can even be used in a way to construct GR \cite{HOJMAN197688}. Since evolution of any phase space variables representing a given physical field as e.g. $(q_{ab},P^{ab})$ for gravitational degrees of freedom (but this applies to any fields), is no more than transporting this variables from one hypersurface to another, the functional which generates evolution should be representations of the generators of the hypersurface deformation. Obviously, the algebra of hypersurface deformation depends on the signature of the original 4-dimensional manifold in which hypersurfaces are embedded. And this is why the algebra of constraints, being a representation of the hypersurface deformation, traces the signature of the original 4-dimensional structure. Following this interpretation and the constructive approach of \cite{HOJMAN197688}, the algebra of constraints, Eqs. (\ref{eq:algebra1}) to (\ref{eq:algebra3}), is a consequence of the algebra of the hypersurface deformation. \\ Over the last decade, results in loop quantum cosmology (LQC hereafter) \cite{Barrau:2013ula}, also in the case of spherically symmetric space-times or Gowdy space-times \cite{Bojowald:2015zha,Bojowald:2015sta}, surprisingly show that at an effective level, the algebra of {\it quantum-corrected constraints} might be deformed as compared to the one obtained in GR \cite{eucl3}. More specifically in the so-called {\it deformed-algebra} approach of LQC, it can be shown that to preserve an algebraic structure, the last Poisson bracket should be modified to $\left\{S(M),S(N)\right\}=D\left(\OM q^{ab}(M\partial_bN-N\partial_bM)\right)$ where $\OM$ is a function of the gravitational phase-space variables. Remarkably enough, this function can even change its sign in some concrete cases: it is positive-valued well below the Planck scale but becomes {\it negative-valued} close to the Planck scale. If the sign in front of the left-hand-side of the Poisson bracket in Eq. (\ref{eq:algebra3}) is to be interpreted as the signature of the underlying 4-dimensional structure, then this would mean that at the Planck scale, space-time becomes Euclidean, or stated otherwise, that time has disappeared. Obviously the full picture is more subtle than that. In GR, we know the underlying 4-dimensional manifold to be pseudoriemanian, and the equivalence between the algebra of hypersurface deformation and the algebra of constraints is clearly established. In LQC however, the deformation of the last Poisson bracket is obtained at the level of the algebra of constraints. Whether this traces or not a deformation of an underlying, signature-changing, 4-dimensional space in which time may disappear, is still highly debated. (What is only certain is that if such a structure does exist, this cannot be a pseudoriemanian space.) More specifically in LQC, the presence of $\OM$ in the algebra of constraints could be interpreted as a mere instability of fields propagating on a quantum space-time, and not as a true signature change of the quantum-corrected space-time. The good point however is that depending on the adopted interpretation, this leads to different observational consequences in the amplitude of cosmological inhomogeneities produced in the early Universe, which in principle open the window for observational tests of this remarkable feature. \\ This article aims at presenting the possible observational consequences of signature change as obtained in LQC. Loop quantum gravity is a well defined theory (see, {\it e.g.} \cite{lqg3} for a review). It is based on Ashtekar's formulation of GR and can be derived from different approaches: canonical quantization of GR, covariant quantization of GR, and quantization of geometry. There are interesting and promising attempts to derive the cosmological dynamics from the full theory (see, {\it e.g.}, \cite{Bianchi:2010zs}). However most of the results obtained in LQC, in particular when dealing with perturbations, are based either on effective equations or on heavy hypotheses. A possible way to address this question is to try to deal with quantum fields on a quantum background, as done in the {\it dressed metric} approach (see, {\it e.g.}, \cite{agullo2}). Another way to face the problem of perturbations is to put the emphasis on the consistency of the effective equations, as done in the {\it deformed-algebra} approach (see, {\it e.g.}, \cite{eucl3}). Here, we mainly -- but not only -- focus on the {\it deformed-algebra} approach which can also be seen as a kind of embedding of GR in a more general framework. In the first part, we recall the reasons why a change of signature can be expected in LQC. In the second section, we explain the consequences of this phenomenon if it is understood as an effective instability in the equation of motion of modes in Fourier space. In the third section, we focus on the consequences if the phenomenon is understood as a real change of signature at the deepest level. Then, we explore in the next section some possible consequences for black holes. Finally, we briefly discuss some open questions related to the signature change in the last section, and which might be of phenomenological relevance.
\label{sec:conclu} The apparent change of signature obviously raises some questions, since unlike GR, it is still unclear whether the deformed algebra of constraints really traces a deformation of the algebra of some "hypersurfaces" embedded in a 4-dimensional, signature-changing space. In GR, there is a clear equivalence between the Hamiltonian formulation and the geometrical formulation, while the formalism presented above is solely phrased in an Hamiltonian language. One may however try to reconstruct some notions such as 4-dimensional, {\it effective} world-lines as a way to have some clues on what would be a possible underlying 4-dimensional structure leading to signature change. Physically speaking, a geodesic is no more than the world-line followed by some matter content in a given space-time. From a field theoretic viewpoint, geodesics are obtained by taking the optical (or eikonal) limit of the fields equation. By performing this, one can easily recover the fact that photons follow null geodesics from the optical limit of the Maxwell equations in curved space-time \cite{Fleury:2015hgz}, or similarly by starting from the Klein-Gordon equation for a massive scalar field, that massive particles follow time-like geodesics (as a result of the optical limit). If the deformation of the algebra of constraints truly traces a change of signature, this should affect any test fields propagating in the quantum-corrected space-time. An interesting idea would thus be to use the optical limit of test fields propagating in the quantum-corrected space as a way to reconstruct the ``world-lines" followed by massless or massive particles. The starting point being the Hamiltonian framework, the strategy could be the following. Considering e.g. a massive scalar field, or the electromagnetic field, the Hamiltonian and diffeomorphism constraints have to be constructed such as to satisfy the {\it deformed} algebra of constraints, thus ensuring that the dynamics of these test fields is generated by proper representations of the algebra of hypersurface deformation (following the idea that the algebra of constraints of any fields has to represent of the algebra of hypersurface deformations \cite{HOJMAN197688}). From that, one thus obtains the deformed equation of motion for the test fields (in the form of a wave equation), from which one finally implements the optical limit to get the properties of the "world-lines" followed by massive particles, or photons. This approach to reconstruct some notions of world-lines is based on the idea that geometrical notions are traced back to the dynamics of some fields, which could be instructive for the underlying 4-dimensional structure (in a very same way that in solid state physics, light propagating in special medium can be described by propagation in a Finsler geometry \cite{Skakala:2008jp}). It may however not be unique. As an exemple, one may argue that the characteristics of the wave equations, as studied in \cite{Bojowald:2015gra}, is another way to derive some notion of effective world-lines. This approach could be interesting not only to shed light on the possible 4-dimensional structure underlying the deformed algebra of constraints, but also in order to search for phenomenological consequences of the change of signature. It is expected from the above procedure to have an equation of motion for point-like particles, which however propagate in a signature-changing space. This could for example serve for studying the behavior of particles traversing e.g. bouncing black holes.
16
7
1607.07589
Within some approaches to loop quantum cosmology, the existence of an Euclidean phase at high density has been suggested. In this article, we try to explain clearly what are the observable consequences of this possible disappearance of time. Depending on whether it is a real fundamental effect or just an instability in the equation of motion, we show that very different conclusions should be drawn. We finally mention some possible consequences of this phenomenon in the black hole sector.
false
[ "high density", "time", "quantum cosmology", "the black hole sector", "Euclidean", "motion", "very different conclusions", "quantum", "some possible consequences", "an Euclidean phase", "this possible disappearance", "the observable consequences", "the equation", "the existence", "this phenomenon", "a real fundamental effect", "just an instability", "some approaches", "this article", "We" ]
10.567753
-0.048108
89
12435852
[ "Kian Tan, Peng", "Kurtsiefer, Christian" ]
2016arXiv160705897K
[ "Characterization of very narrow spectral lines with temporal intensity interferometry" ]
1
[ "-", "-" ]
[ "2017AJ....153..251T" ]
[ "astronomy", "physics" ]
4
[ "Astrophysics - Instrumentation and Methods for Astrophysics", "Quantum Physics" ]
[ "1958PhRv..112.1940S", "1958RSPSA.243..291B", "1960Natur.187..493M", "1961PhRvL...6..106J", "1963PhRv..131.2766G", "1964AmJPh..32..919M", "1965Natur.208...29W", "1965PhRvL..15..912A", "1966PhRvL..17..663S", "1968PhRv..175.1661S", "1970sfss.coll..134M", "1971JOSA...61.1301E", "1974iiia.book.....H", "1986Ap&SS.125..341V", "1995ocqo.book.....M", "2005A&A...435.1181S", "2005NewA...10..361J", "2010ApJ...708..158G", "2012MNRAS.427..581R", "2014ApJ...789L..10T", "2014SPIE.9146E..0ZD", "2015NatCo...6.6852D", "2016MNRAS.457.4291T" ]
[ "10.48550/arXiv.1607.05897" ]
1607
1607.05897_arXiv.txt
Time-resolved second order correlation spectroscopy was used to identify the presence of very narrow-band light on a thermal background. The linewidth of pseudo-thermal light could be determined that was generated by phase-randomization in a multiple scattering process, similar to light from an ensemble of emitters without a fixed phase relationship, like a gas cloud excited by a nearby star. Temporal intensity interferometry offers a spectral resolution of at least a few 10\,MHz for emission lines, exceeding by far that of contemporary astrophysical spectrographs \citep{griest:10}. Also, an identification of sub-thermal photon statistics can be carried out with the presented technique indicating a possible optical lasing mechanism, and therefore help to better understand the very narrow spectral features of stellar light sources even in presence of a strong blackbody radiation background.
16
7
1607.05897
Context: Some stellar objects exhibit very narrow spectral lines in the visible range additional to their blackbody radiation. Natural lasing has been suggested as a mechanism to explain narrow lines in Wolf-Rayet stars. However, the spectral resolution of conventional astronomical spectrographs is still about two orders of magnitude too low to test this hypothesis. Aims: We want to resolve the linewidth of narrow spectral emissions in starlight. Methods: A combination of spectral filtering with single-photon-level temporal correlation measurements breaks the resolution limit of wavelength-dispersing spectrographs by moving the linewidth measurement into the time domain. Results: We demonstrate in a laboratory experiment that temporal intensity interferometry can determine a 20 MHz wide linewidth of Doppler-broadened laser light, and identify a coherent laser light contribution in a blackbody radiation background.
false
[ "narrow spectral emissions", "narrow lines", "conventional astronomical spectrographs", "a blackbody radiation background", "temporal intensity interferometry", "spectral filtering", "a coherent laser light contribution", "Doppler-broadened laser light", "their blackbody radiation", "starlight", "very narrow spectral lines", "Doppler", "wavelength-dispersing spectrographs", "magnitude", "single-photon-level temporal correlation measurements", "Wolf-Rayet stars", "the time domain", "the linewidth measurement", "the spectral resolution", "the resolution limit" ]
13.789796
10.542931
11
12520974
[ "Petrovich, Cristobal", "Muñoz, Diego J." ]
2017ApJ...834..116P
[ "Planetary Engulfment as a Trigger for White Dwarf Pollution" ]
77
[ "Canadian Institute for Theoretical Astrophysics, University of Toronto, 60 St George Street, ON M5S 3H8, Canada; Centre for Planetary Sciences, Department of Physical &amp; Environmental Sciences, University of Toronto at Scarborough, Toronto, Ontario M1C 1A4, Canada;", "Cornell Center for Astrophysics and Planetary Science, Department of Astronomy, Cornell University, Ithaca, NY 14853, USA" ]
[ "2016ApJ...832..160M", "2016MNRAS.462L..84H", "2016MNRAS.463.2958V", "2017A&A...605A..23P", "2017ApJ...842...67M", "2017ApJ...844..116K", "2017ApJ...844L..16S", "2017ApJ...846..146P", "2017ApJ...849....8M", "2017ApJ...850...50K", "2017MNRAS.465.1008V", "2017MNRAS.465.1499V", "2017MNRAS.465.2053V", "2017MNRAS.468.1575B", "2018ApJ...861...35R", "2018MNRAS.473.2871V", "2018MNRAS.476.3939M", "2018MNRAS.477...93H", "2018MNRAS.479.2649K", "2018MNRAS.479.3814H", "2018MNRAS.480...57S", "2018MNRAS.481.2180V", "2018exha.book.....P", "2019A&A...628A.126M", "2019A&A...632A.113D", "2019ApJ...878...58S", "2019ApJ...886..127M", "2019MNRAS.486.3831V", "2019MNRAS.487..133W", "2019MNRAS.488..153V", "2019MNRAS.489..168G", "2019MNRAS.489.2941V", "2019MNRAS.489.5119M", "2020A&A...643A..34O", "2020ApJ...889...45S", "2020ApJ...904L...3M", "2020MNRAS.492.2437V", "2020MNRAS.492.5561M", "2020MNRAS.493..698M", "2020MNRAS.493.5062V", "2020MNRAS.494..442V", "2020MNRAS.494.2861R", "2020MNRAS.496.2292V", "2020MNRAS.497.5171Z", "2020MNRAS.498.4005O", "2021ApJ...911...50U", "2021ApJ...922....4S", "2021MNRAS.500.3481H", "2021MNRAS.501..507O", "2021MNRAS.501.3806M", "2021MNRAS.502.4479H", "2021MNRAS.504.2853H", "2021MNRAS.504.3375S", "2021MNRAS.505.1557V", "2021MNRAS.506..432S", "2021MNRAS.506.1148V", "2021MNRAS.508.5671L", "2021NatAs...5..451H", "2021orel.bookE...1V", "2022ApJ...924...61L", "2022ApJ...925..178H", "2022ApJ...936...30T", "2022ApJS..259...25H", "2022MNRAS.510.1059B", "2022MNRAS.513.4178O", "2023A&A...674A..52J", "2023ApJ...955L..14S", "2023ApJ...958..120K", "2023MNRAS.518..636F", "2023MNRAS.518.4537V", "2023MNRAS.524.6181O", "2024A&A...686A.123L", "2024ApJ...963..113V", "2024ApJ...966L...4A", "2024MNRAS.529.2910N", "2024MNRAS.530.3302C", "2024arXiv240509399M" ]
[ "astronomy" ]
11
[ "planets and satellites: dynamical evolution and stability", "white dwarfs", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1962AJ.....67..591K", "1962P&SS....9..719L", "1963Icar....2..440H", "1969stph.book.....L", "1979ApJ...231..826F", "1993Icar..106..247G", "1995ApJS...99..189C", "1997Natur.386..254H", "2000MNRAS.315..543H", "2002ApJ...572..556D", "2003ApJ...596..477Z", "2005ApJS..161..394F", "2006A&A...448..641D", "2007A&A...474...77K", "2007ApJ...658..569W", "2007ApJ...671..872Z", "2009A&A...498..517K", "2009A&A...505.1221V", "2009AJ....137.3706T", "2010ApJ...722..725Z", "2010ApJS..190....1R", "2010MNRAS.401..867H", "2010MNRAS.409.1631B", "2010PASP..122..905J", "2011A&A...533A...7B", "2011ApJ...735..109W", "2011ApJ...739...31L", "2011ApJS..197...38D", "2011MNRAS.414..930B", "2011MNRAS.417.2104V", "2012A&A...537A.128R", "2012A&A...538A..74P", "2012ApJ...747..148D", "2012ApJ...753...91K", "2012ApJ...754L..36N", "2012ApJ...761..121M", "2012ApJS..200....3B", "2012MNRAS.424..333G", "2012MNRAS.426.2500F", "2013AJ....145...54T", "2013ApJ...765L...8R", "2013MNRAS.430.1171D", "2013MNRAS.431.1686V", "2013Sci...342..218F", "2014A&A...566A..34K", "2014AJ....147..142C", "2014AREPS..42...45J", "2014AmJPh..82..769T", "2014ApJ...785..116L", "2014ApJ...791L..27Z", "2014ApJ...794....3V", "2014MNRAS.437.1216D", "2014MNRAS.437.1404M", "2014MNRAS.439.2442F", "2014MNRAS.439.3371W", "2014MNRAS.445.2244V", "2015ARA&A..53..409W", "2015ApJ...799...27P", "2015ApJ...799..120B", "2015ApJ...805...75P", "2015ApJ...805..100T", "2015MNRAS.446.1424R", "2015MNRAS.447..747L", "2015MNRAS.447.1049V", "2015MNRAS.449.4221H", "2015MNRAS.451.3453V", "2015MNRAS.451L...1G", "2015MNRAS.452.3610A", "2015Natur.526..546V", "2015PNAS..112.9264M", "2016A&A...589L...6A", "2016ARA&A..54..441N", "2016ApJ...816L..22X", "2016ApJ...818L...7G", "2016MNRAS.456.2070T", "2016MNRAS.456.3671A", "2016MNRAS.458.3060L", "2016MNRAS.458.3904R", "2016MNRAS.460.1086M", "2016MNRAS.462L..84H", "2016NewAR..71....9F", "2016RSOS....350571V", "2016arXiv160601236O", "2017MNRAS.465L..44K" ]
[ "10.3847/1538-4357/834/2/116", "10.48550/arXiv.1607.04891" ]
1607
1607.04891_arXiv.txt
\label{sec:intro} Atmospheric metals are not expected to be present in isolated white dwarfs (WDs) with effective temperatures below $\sim25,000$ K. At these temperatures, radiative forces become too weak \citep{chayer95} to significantly counteract the quick gravitational settling that sinks material heavier than helium in extremely short timescales compared to the typical cooling ages of WDs \citep{FM79,K09}. However, it has been found that $\sim25\%-50\%$ of all field WDs exhibit spectral lines that are indicative of the presence of metals in their atmospheres \citep{Z03,Z10,K14}. The high-metallicity material found in the atmospheres of most of these ``polluted'' WDs is consistent with the composition of rock-forming material \citep{Z07,G12,farihi13,JY14}. This observation suggesting that pollution comes from minor rocky bodies (e.g., asteroids). One possibility is that these rocky bodies get very close to the WD so they can be tidally disrupted and then accreted. Further support of this picture comes from observations of circumstellar disks --revealed by infrared excess in the stellar spectrum-- around many polluted WDs (see \citealt{farihi16} for a recent review). These disks orbit within ${\sim}1R_\odot$, roughly the distance at which the material would reside after the tidal disruption (the Roche radius). All the WDs with detected disks have atmospheric pollution. More recently, this picture has been reinforced by the recent observation of minor bodies transiting the polluted WD 1145+017 \citep{vanderburg15,alonso16,G16,rap16,xu16}. Although the leading explanation for WDs pollution --the accretion of tidally disrupted asteroids-- seems robust and well supported by observations, the underlying dynamical mechanism responsible for placing these rocky bodies in star grazing orbits remains much less constrained and understood. A better understanding of this mechanism can lead to new insights into initial conditions leading to WD pollution, as well as into the long-term dynamics and evolution of the planetary systems around WDs and/or their progenitors (typically A and F stars; see \citealt{veras16} for a recent review on this subject). A theoretical model to explain the WD pollution from planetary dynamical instabilities was put forward by \citet{DS02}. According to their model, a planetary system that is marginally stable throughout the main sequence can become unstable due to stellar mass loss during post-MS evolution. This global instability can then promote some asteroids into star grazing orbits. This idea has been explored in more detail using realistic numerical $N$-body integrations of multi-planet systems (no asteroids) and stellar evolution \citep{veras13,MVV14,VG15}. Similarly, the mass loss of the host star can widen the region around mean-motion resonances where chaotic diffusion of asteroids acts efficiently, leading to their posterior tidal disruption \citep{B11,DWS12,FH14}. As well, mass loss in close binary systems can drive the outermost planetesimals into the chaotic orbits, with one of the possible outcomes being collisions with either one of the stars \citep{kratter12}. Thus far, these proposed dynamical mechanisms rely on generally short-timescale instabilities (either scattering or mean-motion-resonance overlap) triggered (or enhanced) by mass loss or simply by the aging of the planetary systems, and still face some difficulties. In particular, these mechanisms are subject to the following constraints: \begin{enumerate} \item {\it the delivery of material must happen for WDs of all ages.} The observations seem to show that neither the rate of polluted WDs, nor that the level of pollution decreases with the WD cooling age \citep{KGF14,wyatt14}. Thus, to explain the observed pollution rate, the underlying mechanism should be able to deliver enough material into the WD's atmosphere independently of how much time it has passed since the stellar mass loss phase. \item {\it The supply of material into white dwarf grazing orbits must be a steady process.} Both the large observed rate of polluted WDs and the short timescales that follow a disruption event (or order the orbital timescale) require of a sustained process to deliver bodies toward disruption. The formation of a debris disks following disruption can extend the duration of the delivery toward the stellar atmosphere, but its associated timescale is still short compared to the cooling ages of most polluted WDs \citep{veras14b,veras15}. \item {\it The reservoir of rocky material has to be long-lived.} The amount of material waiting to be delivered toward the star cannot be arbitrarily large. A planetesimal disk can be destroyed by a collisional cascade, shattering the rocky bodies down to dust, which can be blown out during the RG and AGB phases by radiation pressure (e.g., \citealt{BW10}). All else being equal, disks with lower surface densities and at larger separations can survive for longer timescales, possibly avoiding this fate (e.g., \citealt{wyatt07,HT10,BW10}). \end{enumerate} In this paper, we propose a new mechanism that overcomes (or at least alleviates) these difficulties. We propose that the nature of the instabilities, which drives the material in a planetesimal disk into disrupting orbits, is secular (not scattering nor driven by mean-motion resonances) and that the instabilities are initiated only at the very end of the the stellar evolution (AGB phase) once a stabilizing, pre-existing planetary system is engulfed by an extended stellar envelope. This mechanism can provide steady pollution over all ages of the WD (overcoming the difficulties 1 and 2), while working for a low surface density disk that remains dynamically cold during the main sequence, and that gets gradually depleted long after mass loss has taken place (addressing difficulty 3). We illustrate how the instabilities arise due to the Kozai-Lidov (KL) mechanism in wide ($\gtrsim100$ AU) stellar binaries, although our proposal is more general and sub-stellar companions and other sources of secular excitation are allowed. We expect that for these wide binaries the possible WD pollution associated with post-AGB dust disks and stellar winds might be negligible (e.g., \citealt{ruyter06,VW09,bilik12,clayton14}).
We have studied a new mechanism to explain the observed metal pollution in white dwarfs through the tidal disruption of small rocky bodies in a planetesimal disk. We propose that one or several planets can shield a planetesimal disk against the KL mechanism due a distant binary companion. Once the host star evolves off the main sequence to become a WD, these planets can be engulfed (most likely during the AGB phase), thus triggering the KL mechanism, and leading to the tidal disruption of the rocky bodies in the planetesimal disk. We have shown that this mechanism can account for the observed accretion rates for WDs with all cooling ages provided that the disks have masses $\sim10^{-4}-10^{-2}M_\oplus$. Our model allows for planetesimal disks with large radial extents, and as a consequence, it presents the following advantages compared to other models: \begin{itemize} \item it provides a steady supply of material (each part of the disk has a different and long disruption timescale), enhancing the probability of observing the pollution of WD atmospheres; \item it allows for low-density surface disks, which can survive internal disruptive collisions over long timescales. \end{itemize} This mechanism is only triggered after the host star has left the main sequence, providing a self-consistent explanation as to why the KL mechanism does not act on the planetesimal disk for the prior few Gyrs. Our estimates indicate that this model can account for a significant fraction of the polluted WDs. Complete searches for companions of WDs might shed light on the significance of our proposal.
16
7
1607.04891
The presence of a planetary system can shield a planetesimal disk from the secular gravitational perturbations due to distant outer massive objects (planets or stellar companions). As the host star evolves off the main sequence to become a white dwarf, these planets can be engulfed during the giant phase, triggering secular instabilities and leading to the tidal disruptions of small rocky bodies. These disrupted bodies can feed the white dwarfs with rocky material and possibly explain the high-metallicity material in their atmospheres. We illustrate how this mechanism can operate when the gravitational perturbations are due to the KL mechanism from a stellar binary companion, a process that is activated only after the planet has been removed/engulfed. We show that this mechanism can explain the observed accretion rates if: (1) the planetary engulfment happens rapidly compared to the secular timescale, which is generally the case for wide binaries (&gt; 100 au) and planetary engulfment during the asymptotic giant branch; (2) the planetesimal disk has a total mass of ∼ {10}<SUP>-4</SUP>-{10}<SUP>-2</SUP>{M}<SUB>\oplus </SUB>. We show that this new mechanism can provide a steady supply of material throughout the entire life of the white dwarfs for all cooling ages and can account for a large fraction (up to nearly half) of the observed polluted white dwarfs.
false
[ "rocky material", "small rocky bodies", "secular instabilities", "planets", "planetary engulfment", "stellar companions", "material", "the observed polluted white dwarfs", "distant outer massive objects", "10}<SUP>-4</SUP>-{10}<SUP>-2</SUP>{M}<SUB>\\oplus", "a white dwarf", "the white dwarfs", "wide binaries", "the secular gravitational perturbations", "the asymptotic giant branch", "gt", "the giant phase", "a stellar binary companion", "the secular timescale", "the planetary engulfment" ]
5.582599
12.138557
75
12522578
[ "Sadoun, Raphael", "Zheng, Zheng", "Miralda-Escudé, Jordi" ]
2017ApJ...839...44S
[ "On the Decreasing Fraction of Strong Lyα Emitters around z ∼ 6-7" ]
36
[ "Department of Physics and Astronomy, University of Utah, 115 South 1400 East, Salt Lake City, UT 84112, USA", "Department of Physics and Astronomy, University of Utah, 115 South 1400 East, Salt Lake City, UT 84112, USA", "Institució Catalana de Recerca i Estudis Avançats, Barcelona, Catalonia ; Institut de Ciències del Cosmos, Universitat de Barcelona (IEEC-UB), Barcelona, Catalonia" ]
[ "2017ApJ...841...19M", "2017ApJ...847...63O", "2018A&A...619A.136M", "2018ApJ...863...92B", "2018MNRAS.477.5406Y", "2018MNRAS.479.2564W", "2018MNRAS.479.4566M", "2018MNRAS.480.5140K", "2019ApJ...872..101B", "2019MNRAS.483.5301G", "2019MNRAS.484.4601S", "2019MNRAS.485.1350W", "2019MNRAS.486.2197B", "2019PASP..131g4101S", "2020ApJ...888....6K", "2020IAUS..341..299S", "2020MNRAS.491.1736K", "2020MNRAS.492.1778M", "2020MNRAS.494..703W", "2021A&A...646L..10P", "2021ApJ...908..219H", "2021ApJ...922..263P", "2021MNRAS.501.5294G", "2021MNRAS.504.1902G", "2021MNRAS.506.2390Q", "2021MNRAS.508.3697G", "2022ApJ...931..126P", "2022JApA...43..104M", "2022MNRAS.512.3243S", "2022arXiv221009612S", "2023A&A...677A..88B", "2023arXiv230503042H", "2024MNRAS.tmp.1509K", "2024arXiv240406548A", "2024arXiv240406569T", "2024arXiv240612070H" ]
[ "astronomy" ]
4
[ "dark ages", "reionization", "first stars", "galaxies: high-redshift", "intergalactic medium", "methods: analytical", "radiative transfer", "Astrophysics - Astrophysics of Galaxies" ]
[ "1965ApJ...142.1633G", "1967ApJ...147..868P", "1974ApJ...187..425P", "1990ApJ...350..216N", "1991ApJ...379..440B", "1997ApJ...490..493N", "1998ApJ...495...80B", "1998ApJ...497...21M", "1998ApJ...501...15M", "1998ApJ...502...59T", "2000ApJ...530....1M", "2000ApJ...535..530G", "2002AJ....123.1247F", "2002ApJ...568L..71Z", "2002ApJ...578...33Z", "2003A&A...397..527S", "2003AJ....126....1W", "2003ApJ...588...65S", "2003Sci...300.1904M", "2004MNRAS.347...59B", "2006AJ....132..117F", "2006ApJ...649...14D", "2007ApJ...662...72B", "2007MNRAS.377.1175D", "2007RPPh...70..627B", "2010ApJ...716..574Z", "2010ApJ...723..869O", "2010ApJ...725L.205F", "2010MNRAS.408.1628S", "2011ApJ...728L...2S", "2011ApJ...743..132P", "2011MNRAS.412.2543C", "2012ApJ...744...83O", "2012ApJ...744..179S", "2012ApJ...760..128M", "2012MNRAS.421.2568D", "2012MNRAS.422.1425C", "2012MNRAS.424.1672D", "2013ApJ...775L..29T", "2013ApJS..208...19H", "2013MNRAS.429.1695B", "2013MNRAS.436.1023B", "2014A&A...562A..52D", "2014ApJ...788...74S", "2014ApJ...794..116Z", "2014ApJ...795...20S", "2014ApJ...797...16K", "2014MNRAS.443.2831C", "2014MNRAS.445.2462M", "2014MNRAS.445.3200S", "2014PASA...31...40D", "2014arXiv1409.4946F", "2015A&A...573A..24C", "2015A&A...578A..83G", "2015ApJ...807..180W", "2015ApJ...813...54T", "2015ApJ...813L...8M", "2015MNRAS.446..566M", "2015MNRAS.447.3402B", "2015MNRAS.453.1843S", "2016A&A...594A..13P", "2016A&A...596A.107P", "2016A&A...596A.108P", "2016ApJ...818...38S", "2016MNRAS.455.1385M", "2016MNRAS.460.1328D", "2018MNRAS.473.1416M" ]
[ "10.3847/1538-4357/aa683b", "10.48550/arXiv.1607.08247" ]
1607
1607.08247_arXiv.txt
Cosmic reionization corresponds to the last major phase transition of the universe during which the intergalactic medium (IGM) transitioned from a highly neutral to a highly ionized state. The detailed history of how reionization proceeded is still poorly constrained owing to the limited number of observational tools to probe neutral gas in the early universe \citep[e.g., reviews by][]{Miralda-Escude03,Barkana.Loeb07,Ferrara.Pandolfi14}. The Gunn-Peterson troughs \citep{Gunn65} observed blueward of the Ly$\alpha$ transition in the spectra of high-redshift quasars indicates that reionization was largely completed by $z\sim 6$ \citep[e.g.,][]{Fan.etal02,White.etal03,Fan.etal06,Becker.etal07,Becker.etal15}. Constraints on reionization at $z \ga 6$ can be obtained from measurements of the Thomson scattering optical depth, $\tau_{\rm ts}$, using cosmic microwave background (CMB) data. As $\tau_{\rm ts}$ depends on the number density of free electrons integrated along the line-of-sight, it can be used to infer a characteristic reionization redshift $z_{\rm reion}$ but is insensitive to the precise reionization history. The latest results from the polarization data of the Planck satellite's Low Frequency Instrument and CMB lensing indicate a value $\tau_{\rm ts} = 0.066 \pm 0.013$ \citep{Planck15}, and those from the High Frequency Instrument give $\tau_{\rm ts} = 0.055 \pm 0.009$ \citep{Planck16a}, corresponding to a mean reionization redshift of $z_{\rm reion} \sim$ 7.8--8.8 \citep{Planck16b}. Ly$\alpha$ emitting galaxies, or Ly$\alpha$ emitters\footnote{\lya emitters usually only refer to galaxies \emph{selected} to have strong \lya emission in narrow-band surveys. In this paper, we use the term \lya emitters more generally for any galaxy with detectable \lya emission} (LAEs) constitute a promising alternative for studying neutral gas at early times and can provide important constraints on the late stages of reionization at $z \sim 6-7$ \citep[see e.g., a recent review by][]{Dijkstra14}. LAEs are young star-forming galaxies in which most of the ionizing photons emitted from hot stars are converted to Ly$\alpha$ photons after recombinations in the interstellar medium (ISM), resulting in strong Ly$\alpha$ emission. As such, they have been predicted to be primary targets in the search for high-redshift galaxies \citep{Partridge.Peebles67}. After Ly$\alpha$ photons escape the interstellar medium around the young stars where they are produced, they experience resonant scattering by neutral hydrogen atoms in the surrounding IGM. The Ly$\alpha$ line emitted from these galaxies therefore contains information on the state of the neutral gas in their vicinity. Before reionization was complete, an absorption imprint should be left on the Ly$\alpha$ emission line of LAEs because of the remaining atomic hydrogen in the IGM. The damped absorption wings of IGM regions with a high neutral fraction are expected to substantially suppress the Ly$\alpha$ emission lines of galaxies behind them \citep[e.g.,][]{Miralda.Rees98,Miralda-Escude98}. Recent observations of Ly$\alpha$ emitting galaxies at high redshift have indeed revealed a reduction in the visibility of LAEs between $z\sim 6$ and $7$. Using ultra-deep narrow-band imaging with the Subaru telescope, \citet{Konno.etal14} found a rapid decline in the Ly$\alpha$ luminosity function (LF) of LAEs from $z=6.6$ to $7.3$. Combined with evidence for no evolution in the ultra-violet (UV) continuum LF over the same redshift interval, they concluded that reionization is likely not complete at $z\sim 7$ and that this may explain the sudden decline of the \lya\ LF of LAEs. The Ly$\alpha$ fraction $X_{\rm Ly\alpha}$, defined as the fraction of objects with strong Ly$\alpha$ emission among Lyman Break Galaxies (LBGs), is slowly rising from $z\sim 3$ to $6$ \citep[e.g.,][]{Stark.etal11}, but then decreases suddenly between $z\sim 6$ and $7$ \citep{Fontana.etal10,Stark.etal10,Pentericci.etal11,Ono.etal12,Schenker.etal12,Caruana.etal14,Schmidt.etal16}. Although one can imagine evolution models for intrinsic galaxy properties such as the escape fraction of ionizing photons or the dust content that might explain this drop in observed Ly$\alpha$ emission \citep[e.g.,][]{Dayal.Ferrara12}, the fraction $X_{\rm Ly\alpha}$ should not decline in a synchronized way for all galaxies in the Universe over a narrow time interval (note that $\Delta z = 1$ corresponds to less than $\sim 200\,{\rm Myr}$ at $z\sim 6$, or $\sim$ 20\% of the age of the Universe at that epoch), so this decline has naturally been interpreted as a signature of the increase in the neutral gas fraction in the IGM towards high redshift. The main difficulty with the simple scenario where a smooth IGM at the end of the reionization epoch is causing this drop of LAEs is that, adopting a simple attenuation model for the transmission of the Ly$\alpha$ line through the intervening IGM, the observed drop in $X_{\rm Ly\alpha}$ between $z=6$ and $z=7$ implies a rapid evolution in the volume-averaged neutral fraction $\left<x_{\rm HI}\right>$ of several tens of percent \citep{Pentericci.etal11,Dijkstra14}. This would demand a late and very sudden reionization scenario, implying a surprisingly rapid rise in the emission rate of ionizing photons, and in tension with the Thomson scattering optical depth measurements by the {\it Wilkinson Microwave Anisotropy Probe} (WMAP; $z_{\rm reion} \sim 10.5$ derived in \citealt{Hinshaw.etal13}), although the latest results from Planck \citep{Planck15,Planck16a,Planck16b} have eased the tension with a late reionization. For this reason, alternative scenarios have been proposed to explain the reduction in Ly$\alpha$ flux at $z \sim 6-7$ without requiring large variations in the neutral gas content of the IGM over this redshift interval. \citet{Bolton.Haehnelt13} suggested that the transmission of the Ly$\alpha$ line might be significantly reduced due to the presence of relatively dense and neutral gas absorbers that are self-shielded against ionizing radiation. The size of these absorbers is reduced as the intensity of the ionizing UV radiation background rises. During the late stages of reionization, the mean free path of ionizing UV photons can increase rapidly as the self-shielded absorbers shrink, causing a rapid change in the ionizing background intensity \citep{Miralda-Escude.etal00,Giallongo15,Madau15,Mitra16,Munoz.etal16}. As a consequence, the Ly$\alpha$ transmission through the intervening IGM might be reduced significantly without requiring a large change in $\left<x_{\rm HI}\right>$ from $z=6$ to $z=7$. However, using results from reionization simulations, \citet{Mesinger.etal15} showed that these self-shielded absorbers cannot fully account for the total IGM opacity required to explain the observed drop in $X_{\rm Ly\alpha}$. In general, the resonant nature of the Ly$\alpha$ line makes it difficult to infer the neutral state of the IGM surrounding LAEs directly from the evolution of $X_{\rm Ly\alpha}$, as radiative transfer effects can significantly alter the transmission of the line \citep[e.g.,][]{Zheng.etal10}. In this paper, we aim to further explore the impact of a rapidly evolving ionizing UV background intensity on the visibility of LAEs at $z\sim 6-7$, taking into account radiative transfer effects. Compared to the previous scenarios mentioned above, we focus our investigation on how the distribution of neutral gas surrounding the LAEs \emph{themselves} is affected by the local UV background when taking into account self-shielding effects. Furthermore, we calculate the full radiative transfer of Ly$\alpha$ photons using a Monte-Carlo approach in order to accurately predict the Ly$\alpha$ properties resulting from the transmission through this gas. For this purpose, we use an analytical description to model the gas distribution and kinematics around LAEs at high-redshift. As the radiative transfer of the Ly$\alpha$ line can be significantly modified depending on the gas kinematics, we also explicitly consider inflow of gas onto the host LAE halo. The paper is organized as follows. In Section \ref{sec:model}, we describe our model for the gas distribution around high-$z$ LAEs, and present the details of the self-shielding and Ly$\alpha$ radiative transfer calculations. The Ly$\alpha$ properties of our modeled LAEs as well as the main results on the evolution of the Ly$\alpha$ fraction are presented in Section \ref{sec:results}. Finally, we discuss the implications of our results and conclude in Section \ref{sec:conclusion}.
\label{sec:conclusion} We construct a simple analytical model to describe the density and velocity distribution of the gas around high-redshift LAEs as a function of their host halo mass and redshift. The gas distribution is represented by a spherically symmetric cloud, consisting of the NFW profile with a core region, surrounded by an infall region and the IGM in Hubble expansion farther away from the LAE. Self-shielding on the gas distribution is computed for two values of the external ionizing background intensity, which can increase rapidly as reionization proceeds. Based on detailed \lya radiative transfer calculations, we find that this model is able to account for the observed decrease in the fraction of \lya emitting galaxies in the interval from $z=6$ to $z=7$ for both UV-bright and UV-faint galaxies, if the background intensity drops moderately \revision[(by a factor $\sim 3-10$)] over this redshift interval. The mechanism of this model is that the rapidly growing ionizing background intensity toward $z\sim 6$ leads to a rapid ionization of the infall region surrounding LAE host halos, which greatly reduces the scattering of \lya photons in this region. As a result, compared to the $z\sim 6$ LAEs, the \lya photons in the red peak that are able to escape and can produce the observed \lya emission line are much more spatially extended for $z\sim 7$ LAEs, and the detectable \lya flux within the small central region corresponding to the commonly used observing apertures drops by a factor of a few. This provides a natural explanation for the drop in the fraction of galaxies with strong \lya emission. In this model, a uniform external ionizing background is adopted. During reionization, however, the UV background is expected to be highly inhomogeneous owing to the complex topology of the reionization process and the discreteness of ionizing sources \citep{Miralda-Escude.etal00,Davies.Furlanetto.16}. Inside the regions of the IGM that are highly ionized (i.e., ${\rm H\, II}$ regions), the local UV background can reach higher intensities than that in the neutral IGM far away from ionizing sources. This mainly happens at the early stages of reionization, before ${\rm H}\,{\rm II}$ regions have overlapped. In the post-overlap phase of reionization, when all the low-density IGM is ionized, the spatial fluctuations in the UV background intensity are reduced, as the mean free path of ionizing photons increases rapidly \citep{Gnedin2000}. In our model, we set the IGM environment close to LAEs as a neutral self-shielded gas, whose spatial extent is determined by the intensity of the local UV background, embedded in a large-scale ionized region. This corresponds to the late stage of reionization, which proceeds \emph{outside-in} after the overlap of ${\rm H\, II}$ regions. Therefore, we believe that our adoption of a uniform background is a reasonable approximation around LAEs at $z\sim 6-7$, but of course improved predictions can be achieved with fully three-dimensional radiative transfer cosmological simulations. We have not attempted to model the gas distribution inside the halo virial radius. We have simply introduced a core radius in the neutral gas distribution, ignoring the effects of complex physical processes of shock heating or internal ionization and winds, arguing that the radiative transfer process that matters for the observable \lya emission line at the redshifts of interest occurs in the infall region and not within the virialized halo. We have also assumed a turbulent dispersion and a smooth gas distribution, and have not modeled the effects of gas inflow and outflow, a possible multiphase distribution including clumps \citep[e.g.,][]{Dijkstra12,Duval14}, and an anisotropic gas distribution \citep[e.g.,][]{Zheng.Wallace14}. All these factors can affect the \lya radiative transfer and can produce anisotropic \lya emission, which may modify the \lya EW distribution. Our model focusing on the infall and IGM regions can be regarded as describing an average effect on the transfer of photons escaping the host halo. For future work, a detailed investigation with a more realistic gas distribution can come from modeling LAEs in high-resolution hydrodynamic galaxy formation simulations. \revision[As a test for the uncertainty associated with the gas distribution inside halos, we have performed additional test runs by varying the key physical parameters that affect the neutral gas distribution within $r<r_h$, namely the size of the constant density core and the escape fraction of ionizing photons emitted by the LAE (see Appendix). These tests have shown that even when the neutral gas content inside the halo is much lower than in the fiducial case, our model still predicts a \lya flux decrement from $z\sim 6$ to $7$ at a level similar to that in the fiducial model. This suggests that our main results on the reduced visibility of LAEs towards $z\sim 7$ are robust even if there are uncertainties in our modeling of the neutral gas distribution within the halo, as long as a variations in the UV background intensity are able to affect the ionization state of the gas in the infall region due to self-shielding effects.] \citet{Dijkstra.etal07} also use the infall model of \citet{Barkana04} to study the IGM transmission to \lya emission at $z\ga 6$ by modifying a starting \lya line profile based on the \lya scattering optical depth at each frequency (i.e., the $e^{-\tau}$ model). In our work, after a self-consistent self-shielding correction, we track the scatterings of \lya photons not only inside the halos but also in the infall regions and the IGM. The radiative transfer in the infall and IGM regions leads to additional frequency and spatial diffusion of \lya photons that the $e^{-\tau}$ model cannot capture \citep[e.g.,][]{Zheng.etal10}. While the results may be qualitatively similar, our model treats in greater detail the self-shielding and radiative transfer effects. In our model, the apparent \lya fraction evolution is caused by the changes in neutral gas environment around LAEs induced by the rapid evolution in the UV background intensity. This rapid evolution of the UV background is expected as the mean free path of ionizing photons is quickly rising as a consequence of the reduced number density of optically thick systems towards the end of reionization \citep{Miralda-Escude.etal00}. \citet{Bolton.Haehnelt13} also propose that the Ly$\alpha$ fraction evolution can be explained by the rapid change in the UV background level. Their model differs from ours in that they attribute the reduction of the LAE visibility to the generally increased number density of self-shielded regions from $z\sim 6$ to $7$, whereas we propose that the main effect is due to the change in self-shielding of the infall regions around the host halos of the LAEs themselves. The models may be to some extent complementary and a combination of both reasons may provide a more complete picture on the \lya fraction evolution at $z \ga 6$ \citep[e.g.,][]{Dijkstra14}.
16
7
1607.08247
The fraction of galaxies with strong Lyα emission has been observed to decrease rapidly with redshift at z ≳ 6, after a gradual increase at z &lt; 6. This has been interpreted as being a trace of the reionization of the intergalactic medium (IGM): the emitted Lyα photons would be scattered by an increasingly neutral IGM at z &gt; 6. We study this effect by modeling the ionization and Lyα radiative transfer in the infall region and the IGM around a Lyα emitting galaxy (LAE), for a spherical halo model with the mean density and radial velocity profiles in the standard ΛCDM cosmological scenario. We find that the expected fast increase of the ionizing background intensity toward the end of the reionization epoch implies a rapid evolution of halo infall regions from being self-shielded against the external ionizing background to being mostly ionized. Whereas self-shielded infall regions can scatter the Lyα photons over a much larger area than the commonly used apertures for observing LAEs, the same infalling gas is no longer optically thick to the Lyα emission line after it is ionized by the external background, making the Lyα emission more compact and brighter within the observed apertures. Based on this simple model, we show that the observed drop in the abundance of LAEs at z &gt; 6 does not imply a rapid increase with redshift of the fraction of the whole IGM volume that is atomic, but is accounted for by a rapid increase of the neutral fraction in the infall regions around galaxy host halos.
false
[ "strong Lyα emission", "halo infall regions", "Lyα", "z ≳", "galaxy host halos", "z", "IGM", "the Lyα emission line", "a Lyα emitting galaxy", "the emitted Lyα photons", "gt", "the Lyα emission", "self-shielded infall regions", "redshift", "the external ionizing background", "the ionizing background intensity", "a rapid increase", "the Lyα photons", "the standard ΛCDM cosmological scenario", "galaxies" ]
13.598975
6.22288
159
4924162
[ "Decarli, Roberto", "Walter, Fabian", "Aravena, Manuel", "Carilli, Chris", "Bouwens, Rychard", "da Cunha, Elisabete", "Daddi, Emanuele", "Ivison, R. J.", "Popping, Gergö", "Riechers, Dominik", "Smail, Ian R.", "Swinbank, Mark", "Weiss, Axel", "Anguita, Timo", "Assef, Roberto J.", "Bauer, Franz E.", "Bell, Eric F.", "Bertoldi, Frank", "Chapman, Scott", "Colina, Luis", "Cortes, Paulo C.", "Cox, Pierre", "Dickinson, Mark", "Elbaz, David", "Gónzalez-López, Jorge", "Ibar, Edo", "Infante, Leopoldo", "Hodge, Jacqueline", "Karim, Alex", "Le Fevre, Olivier", "Magnelli, Benjamin", "Neri, Roberto", "Oesch, Pascal", "Ota, Kazuaki", "Rix, Hans-Walter", "Sargent, Mark", "Sheth, Kartik", "van der Wel, Arjen", "van der Werf, Paul", "Wagg, Jeff" ]
2016ApJ...833...69D
[ "ALMA Spectroscopic Survey in the Hubble Ultra Deep Field: CO Luminosity Functions and the Evolution of the Cosmic Density of Molecular Gas" ]
112
[ "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 ; National Radio Astronomy Observatory, Pete V. Domenici Array Science Center, P.O. Box O, Socorro, NM 87801, USA", "Núcleo de Astronomía, Facultad de Ingeniería, Universidad Diego Portales, Av. Ejército 441, Santiago, Chile", "National Radio Astronomy Observatory, Pete V. Domenici Array Science Center, P.O. Box O, Socorro, NM 87801, USA ; Cavendish Laboratory, University of Cambridge, 19 J. J. Thomson Avenue, Cambridge CB3 0HE, UK", "Leiden Observatory, Leiden University, P.O. Box 9513, NL2300 RA Leiden, The Netherlands", "Centre for Astrophysics and Supercomputing, Swinburne University of Technology, Hawthorn, Victoria 3122, Australia ; Research School of Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611, Australia", "Laboratoire AIM, CEA/DSM-CNRS-Université Paris Diderot, Irfu/Service d'Astrophysique, CEA Saclay, Orme des Merisiers, F-91191 Gif-sur-Yvette cedex, France", "European Southern Observatory, Karl-Schwarzschild-Strasse 2, D-85748, Garching, Germany ; Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ, UK", "European Southern Observatory, Karl-Schwarzschild-Strasse 2, D-85748, Garching, Germany", "Cornell University, 220 Space Sciences Building, Ithaca, NY 14853, USA", "6 Centre for Extragalactic Astronomy, Department of Physics, Durham University, South Road, Durham DH1 3LE, UK", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-053121 Bonn, Germany", "Departamento de Ciencias Físicas, Universidad Andres Bello, Fernandez Concha 700, Las Condes, Santiago, Chile", "Departamento de Ciencias Físicas, Universidad Andres Bello, Fernandez Concha 700, Las Condes, Santiago, Chile; Millennium Institute of Astrophysics (MAS), Nuncio Monseñor Sótero Sanz 100, Providencia, Santiago, Chile", "Núcleo de Astronomía, Facultad de Ingeniería, Universidad Diego Portales, Av. Ejército 441, Santiago, Chile", "Millennium Institute of Astrophysics (MAS), Nuncio Monseñor Sótero Sanz 100, Providencia, Santiago, Chile ; Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, Av. Vicuña Mackenna 4860, 782-0436 Macul, Santiago, Chile ; Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO 80301, USA", "Department of Astronomy, University of Michigan, 1085 South University Ave., Ann Arbor, MI 48109, USA", "Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany", "Dalhousie University, Halifax, Nova Scotia, Canada", "ASTRO-UAM, UAM, Unidad Asociada CSIC, Spain", "Joint ALMA Observatory—ESO, Av. Alonso de Córdova, 3104, Santiago, Chile ; National Radio Astronomy Observatory, 520 Edgemont Rd, Charlottesville, VA 22903, USA", "Joint ALMA Observatory—ESO, Av. Alonso de Córdova, 3104, Santiago, Chile", "National Optical Astronomy Observatory", "Laboratoire AIM CEA/DSM-CNRS-Université 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", "Instituto de Física y Astronomía, Universidad de Valparaíso, Avda. Gran Bretaña 1111, Valparaiso, Chile", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, Av. Vicuña Mackenna 4860, 782-0436 Macul, Santiago, Chile", "Leiden Observatory, Leiden University, P.O. Box 9513, NL2300 RA Leiden, The Netherlands", "Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille), UMR 7326, F-13388 Marseille, France", "Argelander Institute for Astronomy, University of Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany", "IRAM, 300 rue de la piscine, F-38406 Saint-Martin d'Hères, France", "Astronomy Department, Yale University, New Haven, CT 06511, USA", "Cavendish Laboratory, University of Cambridge, 19 J. J. Thomson Avenue, Cambridge CB3 0HE, UK ; Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Max-Planck Institut für Astronomie, Königstuhl 17, D-69117, Heidelberg, Germany", "Astronomy Centre, Department of Physics and Astronomy, University of Sussex, Brighton BN1 9QH, UK", "NASA Headquarters, Washington, DC 20546-0001, USA", "Max-Planck Institut für Astronomie, Königstuhl 17, D-69117, Heidelberg, Germany", "Leiden Observatory, Leiden University, P.O. Box 9513, NL2300 RA Leiden, The Netherlands", "SKA Organization, Lower Withington, Macclesfield, Cheshire SK11 9DL, UK" ]
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[ "astronomy" ]
14
[ "galaxies: evolution", "galaxies: formation", "galaxies: high-redshift", "galaxies: ISM", "surveys", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.3847/1538-4357/833/1/69", "10.48550/arXiv.1607.06770" ]
1607
1607.06770_arXiv.txt
The cosmic star-formation history (SFH) describes the evolution of star formation in galaxies across cosmic time. It is well summarized by the so-called ``Lilly-Madau'' plot \citep{lilly95,madau96}, which shows the redshift evolution of the star-formation rate (SFR) density, i.e., the total SFR in galaxies in a comoving volume of the universe. The SFR density increases from an early epoch ($z>8$) up to a peak ($z\sim 2$) and then declines by a factor $\sim 20$ down to present day \citep[see][for a recent review]{madau14}. Three key quantities are likely to drive this evolution: the growth rate of dark matter halos, the gas content of galaxies (i.e., the availability of fuel for star formation), and the efficiency at which gas is transformed into stars. Around $z$=$2$, the mass of halos can grow by a factor of $>2$ in a Gyr; by $z\approx0$, the mass growth rate has dropped by an order of magnitude \citep[e.g.,][]{griffen16}. How does the halo growth rate affect the gas resupply of galaxies? Do galaxies at $z\sim 2$ harbor larger reservoirs of gas? Are they more effective at high redshift in forming stars from their gas reservoirs, possibly as a consequence of different properties of the interstellar medium, or do they typically have more disturbed gas kinematics due to gravitational interactions? To address some of these questions, we need a census of the dense gas stored in galaxies and available to form new stars as a function of cosmic time, i.e., the total mass of gas in galaxies per comoving volume [$\rho$(gas)]. The statistics of Ly$\alpha$ absorbers (associated with atomic hydrogen, H{\sc i}) along the line of sight toward bright background sources provide us with a measure of $\rho$(H{\sc i}). This appears to be consistent with being constant (within a $\sim$30\% fluctuation) from redshift $z=0.3$ to $z\sim 5$ \citep[see, e.g.,][]{crighton15}, possibly as a result of the balance between gas inflows and outflows in low-mass galaxies \citep{lagos14} and of the on-going gas resupply from the intergalactic medium \citep{lagos11}. However, beyond the local universe, little information currently exists on the amount of molecular gas that is stored in galaxies, $\rho$(H$_2$), which is the immediate fuel for star formation (e.g., see review by \citealt{carilli13}). Attempts have been made to infer the mass of molecular gas in distant targeted galaxies indirectly from the measurement of their dust emission, via dust--to--gas scaling relations \citep{magdis11,magdis12,scoville14,scoville15,groves15}. But a more direct route is to derive it from the observations of rotational transitions of $^{12}$CO (hereafter, CO), the second most abundant molecule in the universe (after H$_2$). As the second approach is most demanding in terms of telescope time, it has been traditionally applied only with extreme, infrared (IR) luminous sources (e.g., \citealt{bothwell13}; these however account for only 10-20\% of the total SFR budget in the universe; see, \citealt{rodighiero11,magnelli13,gruppioni13,casey14}), or on samples of galaxies pre-selected based on their stellar mass and/or SFR \citep[e.g.,][]{daddi10a,daddi10b,daddi15,tacconi10,tacconi13,genzel10,genzel15,bolatto15}. These observations have been instrumental in shaping our understanding of the molecular gas properties in high-$z$ galaxies. Through the observation of multiple CO transitions for single galaxies, the CO excitation has been constrained in a variety of systems \citep{weiss07,riechers11,bothwell13,spilker14,daddi15}. Most remarkably, various studies showed that $M_*$- and SFR-selected galaxies at $z>0$ tend to host much larger molecular gas reservoirs than typically observed in local galaxies for a given stellar mass ($M_*$) suggesting that an evolution in the gas fraction $f_{\rm gas}=M_{\rm H2}/(M_*+M_{\rm H2})$ occurs through cosmic time \citep{daddi10a,riechers10,tacconi10,tacconi13,genzel10,genzel15,geach11,magdis12,magnelli12}. For molecular gas observations to constrain $\rho$(H$_2$) as a function of cosmic time, we need to sample the CO luminosity function in various redshift bins. CO is the second most abundant molecule in the universe (after H$_2$) and therefore is an excellent tracer of the molecular phase of the gas. The CO(1-0) ground transition has an excitation temperature of only $T_{\rm ex}=5.5$\,K, i.e., the molecule is excited in virtually any galactic environment. Other low-J CO lines may be of practical interest, as these levels remain significantly excited in star-forming galaxies; and thus, the associated lines [CO(2-1), CO(3-2), CO(4-3)] are typically brighter and easier to detect than the ground state transition CO(1-0). There have been various predictions of the CO luminosity functions both for the J=1$\rightarrow$0 transition and for intermediate and high-J lines, using either theoretical models \citep[e.g.,][]{obreschkow09a,obreschkow09b,lagos11,lagos12,lagos14,popping14a,popping14b,popping16} or empirical relations \citep[e.g.,][]{sargent12,sargent14,dacunha13,vallini16}. Theoretical models typically rely on semi-analytical estimates of the budget of gas in galaxies (e.g., converting H{\sc i} into H$_2$ assuming a pressure-based argument, as in \citealt{blitz06}; via metallicity-based arguments, as in \citealt{gnedin10,gnedin11}; or based on the intensity of the radiation field and the gas properties, as in \citealt{krumholz08,krumholz09}), and inferring the CO luminosity and excitation via radiative transfer models. These models broadly agree on the dependence of $\rho$(H$_2$) on $z$, at least up to $z\sim 2$, but widely differ in the predicted CO luminosity functions, in particular for intermediate and high J transitions, where details on the treatment of the CO excitation become critical. For example, the models by \citet{lagos12} predict that the knee of the CO(4-3) luminosity function lies at $L'\approx 5\times10^8$\,\Kkmspc{} at $z\sim 3.8$, while the models by \citet{popping16} place the knee at a luminosity about 10 times brighter. Such a spread in the predictions highlight the lack of observational constraints to guide the theoretical assumptions. This study aims at providing observational constraints on the CO luminosity functions and cosmic density of molecular gas via the `molecular deep field' approach. We perform a scan over a large range of frequency ($\Delta \nu/\nu\approx25-30$\,\%) in a region of the sky, and ``blindly'' search for molecular gas tracers at any position and redshift. By focusing on a blank field, we avoid the biases due to pre-selection of sources. This method naturally provides us with a well-defined cosmic volume where to search for CO emitters, thus leading to direct constraints on the CO luminosity functions. Our first pilot experiment with the IRAM Plateau de Bure Interferometer \citep[PdBI; see][]{decarli14} led to the first, weak constraints on the CO luminosity functions at $z>0$ \citep{walter14}. The modest sensitivity (compared with the expected knee of the CO luminosity functions) resulted in large Poissonian uncertainties. These can be reduced now, thanks to the Atacama Large Millimeter/Sub-millimeter Array (ALMA). We obtained ALMA Cycle 2 observations to perform two spatially coincident molecular deep fields, at 3mm and 1mm respectively, in a region of the Hubble Ultra Deep Field \citep[UDF,][]{beckwith06}. The data set of our ALMA Spectroscopic Survey (ASPECS) is described in detail in Paper I of this series \citep{walter16}. Compared with the aforementioned PdBI effort, we now reach a factor of 3--4 better sensitivity, which allows us to sample the expected knee of the CO luminosity functions over a large range of transitions. Furthermore, the combination of band 3 and 6 offers us direct constraints on the CO excitation of the observed sources, thus allowing us to infer the corresponding CO(1-0) emission, and therefore $\rho$(H$_2$). The collapsed cube of the 1mm observations also yields one of the deepest dust continuum observations ever obtained \citep[Paper II of this series,][]{aravena16a}, which we can use to compare the $\rho$(H$_2$) estimates based on CO and the $\rho$(gas) estimates based on the dust emission. This paper is organized as follows: In Sec.~\ref{sec_observations} we summarize the observations and the properties of the data set. In Sec.~\ref{sec_analysis} we describe how we derive our constraints on the CO luminosity functions and on $\rho$(H$_2$) and $\rho$(gas). In Sec.~\ref{sec_discussion} we discuss our results. Throughout the paper we assume a standard $\Lambda$CDM cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm m}=0.3$ and $\Omega_{\Lambda}=0.7$ \citep[broadly consistent with the measurements by the][]{planck15}.
\label{sec_discussion} In this paper we use our ALMA molecular scans of the {\em Hubble} UDF in band~3 and band~6 to place blind constraints on the CO luminosity function up to $z\sim 4.5$. We provide constraints on the evolution of the cosmic molecular gas density as a function of redshift. This study is based on galaxies that have been blindly selected through their CO emission, and not through any other multi--wavelength property. The CO number counts have been corrected for by two parameters, {\em fidelity} and {\em completeness}, which take into account the number of false positive detections due to noise peaks and the fraction of lines that our algorithm successfully recovers in our data cubes from a parent population of known (artificial) lines. We start by constructing CO luminosity functions for the respective rotational transitions of CO for both the 3\,mm and 1\,mm observations. We compare these measurements to models that also predict CO luminosities in various rotational transitions, i.e. no assumptions were made in comparing our measurements to the models. This comparison shows that our derived CO luminosity functions lie above the predictions in the 3\,mm band. On the other hand, in the 1\,mm band our measurements are comparable to the models. Together this implies that the observed galaxies are more gas--rich than currently attributed for in the models, but with lower excitation. Accounting for a CO excitation characteristic of main--sequence galaxies at $z\sim 1$--2, we derive the CO luminosity function of the ground--transition of CO (J=1--0) from our observations. We do so only up the J=4 transition of CO, to ensure that our results are not too strongly affected by the excitation corrections that would dominate the analysis at higher J. We find an evolution in the CO(1-0) luminosity function compared with observations in the local universe, with an excess of CO-emitting sources at the bright end of the luminosity functions. This is in general agreement with first constraints on the CO intensity mapping from the literature. This evolution exceeds what is predicted by the current models. This discrepancy appear to be a common trait of models of galaxy formation: galaxies with $M_*>10^{10}$\,\Msun{} at $z=2-3$ are predicted to be 2--3 times less star forming than observed (see, e.g., the recent review by \citealt{somerville15}), and similarly less gas--rich (see the analysis in \citealt{popping15a,popping15b}). The sensitivity of the ALMA observations reaches below the knee of the predicted CO luminosity functions (around 5$\times$10$^{9}$\,\Kkmspc) at all redshifts. We convert our luminosity measurements into molecular gas masses via a `Galactic' conversion factor. By summing the molecular gas masses obtained at each redshift, we obtain an estimate of the cosmic density of molecular gas in galaxies, $\rho$(H$_2$). Given the admittedly large uncertainties (mainly due to Poisson errors), and the unknown shape of the intrinsic CO luminosity functions, we do not extrapolate our measurements outside the range of CO luminosities (i.e., H$_2$ masses) covered in our survey. We find an increase (factor of 3--10) of the cosmic density of molecular gas from $z\sim 0$ to $z\sim$ 2--3, albeit with large uncertainties given the limited statistics. This is consistent with previous findings that the gas mass fraction increases with redshift \citep[see, e.g.,][]{tacconi10,tacconi13,magdis12}. However our measurements have been derived in a completely different fashion, by simply counting the molecular gas that is present in a given cosmic volume, without any prior knowledge of the general galaxy population in the field. In this respect, our constraints on $\rho$(H$_2$) are actually lower limits, in the sense that they do not recover the full extent of the luminosity function. However, a) we do sample the predicted knee of the luminosity function in most of the redshift bins, suggesting that we recover a large part ($>$50\%) of the total CO luminosity per comoving volume; b) the fraction of the CO luminosity function missed because of our sensitivity cut is likely larger at higher redshift, i.e., correcting for the contribution of the faint end would make the evolution in $\rho$(H$_2$) even steeper. We have also derived the molecular gas densities using the dust emission as a tracer for the molecular gas, following \citet{scoville14,scoville15}. The molecular gas densities derived from dust emission are generally smaller than but broadly consistent with those measured from CO at $z<2$, but that they might fall short at reproducing the predicted gas mass content of galaxies at $z>2$. Our analysis demonstrates that CO-based gas mass estimates result in 3--10 times higher gas masses in galaxies at $z\sim2$ than in the local universe. The history of cosmic SFR \citep{madau14} appears to at least partially follow the evolution in molecular gas supply in galaxies. The remaining difference between the evolution of the SFR density (a factor of $\sim 20$) and the one of molecular gas (a factor of 3--10) may due to the shortened depletion time scales. A further contribution to this difference may be ascribed to cosmic variance. The UDF in general (and therefore also the region studied here) is found to be underdense at $z>3$ \citep[e.g., Fig.~14 in][]{beckwith06} and in IR-bright sources \citep{weiss09}. The impact of cosmic variance can be estimated empirically from the comparison with the number counts of sources detected in the dust continuum \citep{aravena16a}, or analytically from the variance in the dark matter structures, coupled with the clustering bias of a given galaxy population \citep[see, e.g.,][]{somerville04}. \citet{trenti08} provide estimates of the cosmic variance as a function of field size, halo occupation fraction, survey completeness, and number of sources in a sample. For a $\Delta z=1$ bin centered at $z=2.5$, a 100\% halo occupation fraction and 5 sources detected over 1 arcmin$^2$ (i.e., roughly mimicing the $z\sim2.5$ bin in our analysis), the fractional uncertainty in the number counts due to cosmic variance is $\sim 20$\% ($\sim 60$\% if we include Poissonian fluctuations). Already a factor 5 increase in target area (resulting in a field that is approximately the size of the {\em Hubble} eXtremely Deep Field, \citealt{illingworth13}), at similar depth, would beat down the uncertainties significantly ($\lsim 30$\,\%, including Poissonian fluctuations). With ALMA now being fully operational, such an increase in areal coverage appears to be within reach.
16
7
1607.06770
In this paper we use ASPECS, the ALMA Spectroscopic Survey in the Hubble Ultra Deep Field in band 3 and band 6, to place blind constraints on the CO luminosity function and the evolution of the cosmic molecular gas density as a function of redshift up to z ∼ 4.5. This study is based on galaxies that have been selected solely through their CO emission and not through any other property. In all of the redshift bins the ASPECS measurements reach the predicted “knee” of the CO luminosity function (around 5 × 10<SUP>9</SUP> K km s<SUP>-1</SUP> pc<SUP>2</SUP>). We find clear evidence of an evolution in the CO luminosity function with respect to z ∼ 0, with more CO-luminous galaxies present at z ∼ 2. The observed galaxies at z ∼ 2 also appear more gas-rich than predicted by recent semi-analytical models. The comoving cosmic molecular gas density within galaxies as a function of redshift shows a drop by a factor of 3-10 from z ∼ 2 to z ∼ 0 (with significant error bars), and possibly a decline at z &gt; 3. This trend is similar to the observed evolution of the cosmic star formation rate density. The latter therefore appears to be at least partly driven by the increased availability of molecular gas reservoirs at the peak of cosmic star formation (z ∼ 2).
false
[ "z", "more CO-luminous galaxies", "recent semi-analytical models", "significant error bars", "∼", "molecular gas reservoirs", "cosmic star formation", "gt", "CO", "K km s", "galaxies", "the CO luminosity function", "The comoving cosmic molecular gas density", "the cosmic molecular gas density", "blind constraints", "the cosmic star formation rate density", "band", "redshift", "pc", "z &gt" ]
13.008361
7.906714
132
12434671
[ "von Essen, C.", "Cellone, S.", "Mallonn, M.", "Tingley, B.", "Marcussen, M." ]
2016arXiv160703680V
[ "Modelling systematics of ground-based transit photometry I. Implications on transit timing variations" ]
5
[ "-", "-", "-", "-", "-" ]
[ "2018A&A...615A..79V", "2018A&A...620A.142A", "2019A&A...628A.115V", "2021A&A...648A..85E", "2023AJ....165....5G" ]
[ "astronomy" ]
2
[ "Astrophysics - Earth and Planetary Astrophysics" ]
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[ "10.48550/arXiv.1607.03680" ]
1607
1607.03680_arXiv.txt
The advent of highly-accurate long-term and space-based observations such as Kepler \citep{Borucki2010} and CoRoT \citep{Baglin2006} marked a new era for exoplanet search. For instance, Kepler light curves already revealed clear signatures of transit timing variations \citep[TTVs; see e.g.,][]{Holman2010,Lissauer2011,Ballard2011,Steffen2013}, a technique that relies on the variations in the timings of transits to detect and characterize planetary systems with members that can be as light as one Earth mass or below \citep{Agol2005,Holman2005}. Despite their indisputable power, these space missions were neither designed to observe the whole sky nor to follow up already known single exoplanetary systems outside their fields of view. At present, this role can only be played by ground-based telescopes located across the globe. To produce reliable TTV studies, optimal ground-based observations of primary transits would require, to begin with, a sufficiently long time baseline, good phase coverage, and deep primary transits. However, TTV studies are carried out under less strict conditions. Literature already reveals how misleading can be ground-based observations when orbital and physical parameters derived from them are being compared to each other. For instance, after analyzing archival data plus two extra consecutive transit observations, \cite{Diaz2008} reported TTVs in OGLE-Tr-111. Later on, 6 additional transit light curves and a new re-analysis of the complete data set revealed no detection of TTVs \citep{Adams2010}. Another system that has been systematically observed during primary transit is WASP-3. Using observations obtained by means of small aperture telescopes, \cite{Maciejewski2010} firstly reported the detection of TTVs in the system. Additionally, after collecting more than 3 years of transit observations \cite{Eibe2012} reported probable variations in the transit duration, instead of the claimed TTVs. However, \cite{Monatalto2012} studied thirty-eight archival light curves in an homogeneous way, and found no significant evidence of TTVs in the system. Also our group encountered problems in identifying TTVs: although \cite{vonEssen2013} reported indications of TTVs in the Qatar-1 system, \cite{Maciejewski2015} and \cite{Mislis2015} did not reproduce them using more precise and extensive data. Even variations in the inclination were wrongly claimed. Before the Kepler team released the first quarters, \cite{Mislis2010} reported a significant variation in the inclination of \mbox{TrES-2}, one planetary system within Kepler's field of view. Afterwards, \cite{Schroeter2012} re-analyzed all the published observations in addition to the Kepler data, finding that while ground-based observations revealed a declining trend in inclination, Kepler data were consistent with no variation at all \citep[see][Figure 2]{Schroeter2012}. Intriguingly, TrES-2b produces one of the largest primary transit depths and the host star is relatively bright, which would make it an easy target to be observed from the ground. Although the TTV technique is a powerful method to detect exoplanets in multiple systems, the systematic disagreement between authors causes critical readers to disbelieve low-amplitude results. Added to this, planet formation theories \citep{Fogg2007,Mandell2007} and highly precise observations \citep{Steffen2012,Steffen2013b} reveal that hot Jupiters are prone to conform single systems instead of multiple ones. It is natural then to ask ourselves if standard techniques used to analyze ground based transit data are reliable enough to produce robust results, or if the technology used to carry out these observations plus the effects introduced by our atmosphere on photometric data are playing a main role against us. These circumstances motivated us to write a code capable to create realistic synthetic light curves affected by systematics commonly present in ground-based observations. The main goal is to study under which conditions can the artificially added TTV signal be retrieved. In this work we present a detailed description of our code, of the noise sources that are injected into the light curves, and we show a rigorous test that quantifies the resemblance between our synthetic light curves and real data. We show how and to which extent can systematics not properly accounted for reproduce TTV signals and quantify their impact over the characterization of the perturbing planet. We finish our work trying to characterize light curve observables that would be associated to reliable mid-transit times.
In this work we analyze whether current techniques used to detrend transit light curves taken from ground-based telescopes are suitable to properly characterize multiplicity in particular transiting systems via the transit timing variation technique. To this end, we simulated primary transit observations caused by a hot Jupiter which orbital and physical configuration mimics a real system, Qatar-1. To these light curves we artificially added a perturbation in their mid-transit times caused by an Earth-sized planet in a 3:2 mean motion resonance. The synthetic data accounts with what we believe are the most significant sources of light curve deformations: environmental variability (airmass, atmospheric extinction, and chaotic variability in the sky conditions during observations) and instrumental variability. We then tested the quality of our light curves, comparing their noise characteristics to the ones present in real data. As already shown by other authors, our results disfavor the use of incomplete light curves to carry out TTV studies. Furthermore, our studies show that it is more likely that systematics not properly accounted for are causing the observed scatter in the O--C diagram rather than a gravitationally bound exoplanet. This, in consequence, produces mass estimates of the perturbing body that are a factor of up to two larger as expected. We also find that transits normalized by a time-dependent low-order polynomial provide more inaccurate and sometimes even inconsistent orbital and physical parameters than the ones derived from a more instrumentally and environmentally-related detrending function, which includes time-dependent variability such as changes of airmass, seeing, centroid position and integrated flat counts. Nonetheless, our results suggest that either both approaches are insufficient to account for systematics, or error estimates on the mid-transit times are being underestimated by current statistical techniques by a factor of up to three. A final inspection of the O--C diagrams and the light curves associated to each O--C point make us conclude that when only light curves with close-to-full transit coverage, good cadence, and large photometric quality are considered to carry out TTV studies, the derived O--C diagrams appear to be more consistent with their expected variability. In a future work we will investigate if the use of Gaussian Process regression can improve the determination of the perturbers characteristics, which would allow us to use low-quality transit photometry for TTV studies.
16
7
1607.03680
The transit timing variation technique (TTV) has been widely used to detect and characterize multiple planetary systems. Due to the observational biases imposed mainly by the photometric conditions and instrumentation and the high signal-to-noise required to produce primary transit observations, ground-based data acquired using small telescopes limit the technique to the follow-up of hot Jupiters. However, space-based missions such as Kepler and CoRoT have already revealed that hot Jupiters are mainly found in single systems. Thus, it is natural to question ourselves if we are properly using the observing time at hand carrying out such follow-ups, or if the use of medium-to-low quality transit light curves, combined with current standard techniques of data analysis, could be playing a main role against exoplanetary search via TTVs. The purpose of this work is to investigate to what extent ground-based observations treated with current modelling techniques are reliable to detect and characterize additional planets in already known planetary systems. To meet this goal, we simulated typical primary transit observations of a hot Jupiter mimicing an existing system, Qatar-1. To resemble ground-based observations we attempt to reproduce, by means of physically and empirically motivated relationships, the effects caused by the Earth's atmosphere and the instrumental setup on the synthetic light curves. Therefore, the synthetic data present different photometric quality and transit coverage. In addition, we introduced a perturbation in the mid-transit times of the hot Jupiter, caused by an Earth-sized planet in a 3:2 mean motion resonance. Analyzing the synthetic light curves produced after certain epochs, we attempt to recover the synthetically added TTV signal by means of usual primary transit fitting techniques, and show how these can recover (or not) the TTV signal.
false
[ "multiple planetary systems", "single systems", "usual primary transit fitting techniques", "primary transit observations", "Qatar-1", "typical primary transit observations", "current standard techniques", "current modelling techniques", "hot Jupiters", "means", "additional planets", "already known planetary systems", "TTV", "The transit timing variation technique", "data analysis", "different photometric quality and transit coverage", "TTVs", "exoplanetary search", "certain epochs", "an existing system" ]
7.015548
14.030751
-1
12516841
[ "Spaleniak, Izabela", "MacLachlan, David G.", "Gris-Sánchez, Itandehui", "Choudhury, Debaditya", "Harris, Robert J.", "Arriola, Alexander", "Allington-Smith, Jeremy R.", "Birks, Timothy A.", "Thomson, Robert R." ]
2016SPIE.9912E..28S
[ "Modal noise characterisation of a hybrid reformatter" ]
3
[ "Heriot-Watt Univ. (United Kingdom)", "Heriot-Watt Univ. (United Kingdom)", "Univ. of Bath (United Kingdom)", "Heriot-Watt Univ. (United Kingdom)", "Durham Univ. (United Kingdom)", "Heriot-Watt Univ. (United Kingdom)", "Durham Univ. (United Kingdom)", "Univ. of Bath (United Kingdom)", "Heriot-Watt Univ. (United Kingdom)" ]
[ "2017OExpr..2518713Y", "2017OExpr..2525546C", "2018MNRAS.478.4881A" ]
[ "astronomy", "physics" ]
5
[ "Physics - Optics", "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "1980JOSA...70..968R", "2001PASP..113..851B", "2005OptL...30.2545L", "2006ApOpt..45..519C", "2006PhDT.......304C", "2009OExpr..17.1988N", "2010OExpr..18.4673N", "2011MNRAS.417..689L", "2011OExpr..19.5698T", "2012OExpr..2013996B", "2013Nanop...2..429L", "2013OExpr..2127197S", "2014ApJ...786...18M", "2014SPIE.9147E..7PJ", "2015ApJ...806...61H" ]
[ "10.1117/12.2232708", "10.48550/arXiv.1607.03363" ]
1607
1607.03363_arXiv.txt
\label{sec:intro} Multimode fibres have been heavily exploited in astronomy to efficiently couple the light from a telescope focus and transmit it to an astronomical instrument, such as a spectrograph. However multimode fibres suffer from an effect called \emph{modal noise} which is caused by the interference between the fibre modes and refers to noise introduced by the variation of the light pattern in the multimode fibre output \cite{Mahadevan2014}. The modal noise effect is observed when \cite{Rawson1980, Mahadevan2014}: (i) the source spectrum is sufficiently narrow (using a narrow bandwidth laser source or in a high-resolution spectrum) forming a speckle pattern at the fibre output, (ii) some form of spatial filtering is present at the output of the fibre, (iii) the fibre length and position are affected by movement, stress or ambient temperature variation or a change in the source wavelength or a shift in the input illumination. When fibres are used in high resolution spectroscopy, we encounter all of these factors: narrowband light, spatial filtering in the spectrograph by e.g. beam truncation, spatially dependent grating efficiency. Finally the optical fibre linking the telescope and the spectrograph is often under movement and affected by ambient temperature variation and stress. The input illumination also can vary during observations due to errors in the telescope pointing as well as the astronomical seeing variation. All of these effects cause the resulting speckle pattern to fluctuate over time and result in an unstable point spread function (PSF) of the spectrograph and variation of the signal intensity. The mode coupling is highly wavelength dependent and if we assume that conditions (ii) and (iii) are constant, then the signal intensity will vary along the wavelength with an oscillating pattern of sub-nm period\cite{Chen2006b}. As soon as the fibre or coupling into the fibre is disturbed, the periodic pattern changes. On a telescope, normally this change occurs in a timescale of $\sim$\,tens of seconds which is less than the usual time between calibration and observation, introducing uncertainties in the spectral measurements. What is more, the light from a calibration source is often delivered by a single-mode fibre to the multimode fibre link, creating a different speckle pattern to the one created by the astronomical source. To mitigate and average out the pattern variation fluctuation, many types of fibre scramblers have been demonstrated and are in use, e.g. fibre agitators \cite{Baudrand2001,Daino1980,Corbett2006,Lemke2011}, fibre of various geometries (e.g. square, octagonal) \cite{Avila2010}, beam homogenizers \cite{Avila2010}, ball lens scramblers \cite{Halverson2015}, integrating sphere with a moving diffuser \cite{Mahadevan2014}. However the use of single-mode fibres in high-precision Doppler spectrographs has become very attractive in recent years: not only can they provide a stable PSF and an absence of modal noise, but they also allow for the use of a smaller spectrograph with the associated cost and complexity benefits. However, a major obstacle in using single-mode fibres on seeing-limited telescopes is the inefficient coupling of light into the fibre. There has been recent work showing an efficient coupling into single-mode fibres using extreme adaptive-optics systems \cite{Jovanovic2014} but most telescopes still operate in the seeing-limited regime and exhibit poor coupling into the single-mode fibre. In 2005 Leon-Saval et al. \cite{Leon-Saval2005} proposed and demonstrated mode splitters, so called \emph{photonic lanterns}\cite{Leon-Saval2013}, devices which can convert a multimode signal into multiple single-mode signals and vice versa. Since then, there has been a lot of work to improve the device efficiency using either a fibre platform \cite{Noordegraaf2009,Noordegraaf2010,Birks2012} or direct-write technique \cite{Thomson2011,Spaleniak2013a}. The output of the photonic lantern is intrinsically single-mode, therefore should in principle be modal-noise free. However, there have been reports of different kinds of modal noise effects in photonic lanterns \cite{Olaya2012a,NickPasp}, indicating that it should investigated in all new types of devices. In this paper we address the issue of modal noise in a hybrid reformatter. The device is consists of a Multicore Fibre (MCF) photonic lantern and a pseudo-slit reformatter. First we present laboratory results of the variation of the device throughput signal along the wavelength. Then we present the results of the near-field stability of the pseudo-slit.
Modal noise in multimode fibres is a limiting aspect of high-resolution Doppler spectroscopy. Photonic lanterns, in principle do not exhibit this effect, however previous studies have shown that it not always the case. In this paper, we performed two types of measurements to address the modal noise effect in a hybrid device -- a component consisting of a MCF photonic lantern and direct-write pseudo-slit reformatter. We did not find any significant and periodic variation of the device throughput over wavelength at high spectral resolution, which means that we did not detect the modal noise effect in the device. We also found that the device is a very good mode scrambler and gentle agitation of MCF produces a homogeneous image at the device output. However we also found that the PSF of the device output is not completely stable, so it varies depending on excited modes and produced speckles. We showed that this variation originates from the fact that the pseudo-slit is not perfectly straight along its length. The observed PSF instability can contribute to errors in the radial velocities measurements. To overcome this problem we propose either improvement of the laboratory environment and inscription setup stability to produce a ``perfectly'' straight pseudo-slit or fabrication of a one dimensional array of discrete single-mode waveguides at the device output.
16
7
1607.03363
This paper reports on the modal noise characterisation of a hybrid reformatter. The device consists of a multicore-fibre photonic lantern and an ultrafast laser-inscribed slit reformatter. It operates around 1550 nm and supports 92 modes. Photonic lanterns transform a multimode signal into an array of single-mode signals, and thus combine the high coupling efficiency of multimode fibres with the diffraction-limited performance of single-mode fibres. This paper presents experimental measurements of the device point spread function properties under different coupling conditions, and its throughput behaviour at high spectral resolution. The device demonstrates excellent scrambling but its point spread function is not completely stable. Mode field diameter and mode bary-centre position at the device output vary as the multicore fibre is agitated due to the fabrication imperfections.
false
[ "multimode fibres", "high spectral resolution", "Mode field diameter", "different coupling conditions", "single-mode fibres", "function properties", "single-mode signals", "Photonic lanterns", "the high coupling efficiency", "the multicore fibre", "an ultrafast laser-inscribed slit reformatter", "excellent scrambling", "a multicore-fibre photonic lantern", "a hybrid reformatter", "bary-centre position", "its point spread function", "the fabrication imperfections", "experimental measurements", "the device point", "the device output" ]
14.046966
10.262471
11
12516849
[ "Ruane, Garreth", "Jewell, Jeffery", "Mawet, Dimitri", "Pueyo, Laurent", "Shaklan, Stuart" ]
2016SPIE.9912E..2LR
[ "Apodized vortex coronagraph designs for segmented aperture telescopes" ]
22
[ "California Institute of Technology (United States)", "Jet Propulsion Lab. (United States)", "California Institute of Technology (United States)", "Space Telescope Science Institute (United States)", "Jet Propulsion Lab. (United States)" ]
[ "2016SPIE.9909E..0SC", "2017AJ....153...43S", "2017AJ....154..240F", "2017SPIE10398E..0FP", "2017SPIE10400...0GM", "2017SPIE10400E..0EC", "2017SPIE10400E..0JR", "2017arXiv171202042R", "2018AJ....155....7M", "2018AJ....155....8M", "2018Geosc...8..442G", "2018JATIS...4a5004R", "2018SPIE10698E..2SR", "2018SPIE10698E..4UR", "2018SPIE10698E..5XZ", "2019ApJ...873L...7A", "2019JATIS...5d5003C", "2020AJ....159...79L", "2020JOSAA..37..629W", "2022JATIS...8a5003L", "2022JATIS...8b9001L", "2024CRPhy..24S.133G" ]
[ "astronomy", "physics" ]
5
[ "Astrophysics - Instrumentation and Methods for Astrophysics", "Physics - Optics" ]
[ "2006SPIE.6272E..0LM", "2009SPIE.7440E..0OM" ]
[ "10.1117/12.2231715", "10.48550/arXiv.1607.06400" ]
1607
1607.06400_arXiv.txt
\label{sec:intro} % The development of extreme adaptive optics and high-contrast imaging techniques for ground-based telescopes (e.g. GPI \cite{GPI2006}, SPHERE \cite{SPHERE2008}, and SCExAO \cite{Martinache2009}) has enabled the detection and characterization of several young, giant exoplanets \cite{Marois2008,Lafreniere2008,Lagrange2009,Macintosh2015}. However, planets within the detection limits of current instruments are relatively rare \cite{Bowler2016}. The next generation of ground-\cite{Macintosh2006,Kasper2010} and space-based \cite{Noecker2016,Feinberg2014,Dalcanton2015} telescopes will be capable of detecting fainter, older, less massive planets at smaller angular separation from their host stars, thereby providing access to planet populations with significantly higher occurrence rates. Additionally, thorough spectral characterization will be possible for many of these targets thanks to rapid technological developments for precise control and calibration of unwanted stellar radiation, including dedicated coronagraphs for diffraction suppression, wavefront control, as well as new observing and post-processing strategies. Detection and characterization of faint planets requires an optical system that isolates the light from the planet from noise associated with starlight. A coronagraph accomplishes this by manipulating the amplitude and phase of the incoming light such that the diffracted starlight is suppressed or removed optically prior to detection. Several elegant coronagraph designs exist that provide various levels of suppression and planet throughput \cite{Kuchner2002,Kasdin2003,Codona2004,Mawet2005,Foo2005,Guyon2005,Soummer2005,Trauger2007,Guyon2010}. The performance and complexity of each depends on the shape of the telescope pupil. Large, segmented apertures present a unique challenge; the coronagraph masks must be designed to account for the diffraction owing to discontinuities in the aperture including the secondary mirror, spider support structures, and gaps between mirror segments. \cite{Mawet2011_improved,Mawet2013_ringapod,Carlotti2014,Ruane2015_SPIE,Ruane2015_LPM,Pueyo2013,Mazoyer2015,Guyon2014,Balasubramanian2015,Trauger2016,Zimmerman2016} (Also see Zimmerman et al. and Guyon et. al., these proceedings). The vortex coronagraph (VC) \cite{Mawet2005,Foo2005} has been demonstrated to provide high sensitivity to planets at small angular separations \cite{Serabyn2010}. However, complicated aperture shapes limit the performance of the conventional VCs \cite{Mawet2010b} and thereby drive the complexity of the optical design \cite{Mawet2011_improved,Mawet2013_ringapod,Carlotti2014,Ruane2015_SPIE,Ruane2015_LPM} and/or requirements for wavefront control \cite{Pueyo2013,Mazoyer2015}. This work overcomes this technical challenge by introducing a gray-scale apodizer to the VC that acts to suppress polychromatic, diffracted starlight at angular separations $<10\lambda/D$ potentially down to the $10^{-10}$ level on a segmented aperture telescope similar to those proposed for a future LUVOIR flagship mission \cite{Dalcanton2015,SCDA}. \newpage
\label{sec:concl} We have presented apodized pupil vortex coronagraphs designed for ground- and space-based telescopes with segmented apertures. in each case, the coronagraph masks are optimized such that the estimated integration time is minimized, in the presence of noise. The design goals in terms of starlight suppression and throughput for ground- and space-based applications depend mostly strongly on the expected speckle noise characteristics. In case of TMT, the coronagraph masks suppress diffracted starlight to the $10^{-9}$ level, assuming no aberrations in the system, which is well below the expected speckle noise level. For space telescopes, we show solutions that push the diffraction down to the $10^{-10}$ level while maintaining sufficient throughput to significantly reduce integration time estimates. The throughput achieved is relatively insensitive to aperture discontinuities. In the case of the space telescopes, the throughput with and without thick spiders only differs by a couple of percent. However, the central obscuration size has a much larger effect. The best performance is expected for telescopes with relatively small secondary mirrors. Although the theoretical solutions are independent of wavelength, manufacturing errors will ultimately limit the performance of these coronagraph designs. Pathways to approach the ideal performance are available thanks to successful demonstrations of broadband vortex phase masks based on liquid crystal polymers \cite{Mawet2009} and sub-wavelength gratings \cite{Delacroix2013} as well as gray-scale ring apodizers for vortex coronagraphs \cite{Mawet2014}. Detailed simulations are underway to study the chromaticity of gray-scale apodizers fabricated using various methods. Outcomes of these studies will inform a second generation of the presented coronagraph designs that incorporate known material properties. The methods employed here may be readily generalized to include optimization of the deformable mirror shapes to achieve broadband starlight suppression, potentially at high throughput. A comprehensive exploration of apodizer solutions and designs that also make use two deformable mirrors will be the topic of an upcoming paper. \newpage \appendix
16
7
1607.06400
Current state-of-the-art high contrast imaging instruments take advantage of a number of elegant coronagraph designs to suppress starlight and image nearby faint objects, such as exoplanets and circumstellar disks. The ideal performance and complexity of the optical systems depends strongly on the shape of the telescope aperture. Unfortunately, large primary mirrors tend to be segmented and have various obstructions, which limit the performance of most conventional coronagraph designs. We present a new family of vortex coronagraphs with numerically-optimized gray-scale apodizers that provide the sensitivity needed to directly image faint exoplanets with large, segmented aperture telescopes, including the Thirty Meter Telescope (TMT) as well as potential next-generation space telescopes.
false
[ "aperture telescopes", "potential next-generation space telescopes", "faint exoplanets", "nearby faint objects", "most conventional coronagraph designs", "elegant coronagraph designs", "circumstellar disks", "TMT", "the Thirty Meter Telescope", "vortex coronagraphs", "exoplanets", "large primary mirrors", "various obstructions", "image", "the telescope", "starlight", "advantage", "complexity", "numerically-optimized gray-scale apodizers", "The ideal performance" ]
14.429084
10.731449
10
12409048
[ "Fraija, N.", "Marinelli, A." ]
2016ApJ...830...81F
[ "Neutrino, γ-Ray, and Cosmic-Ray Fluxes from the Core of the Closest Radio Galaxies" ]
30
[ "Instituto de Astronomía, Universidad Nacional Autónoma de México, Apdo. Postal 70-264, Cd. Universitaria, DF 04510, México", "Dipartimento di Fisica, Universita di Pisa and I.N.F.N., Largo Bruno Pontecorvo, 3, I-56127 Pisa, Italia" ]
[ "2017APh....89...14F", "2017ApJS..232....7F", "2018MNRAS.481.4461F", "2019A&A...623A...2A", "2019ApJS..245...18F", "2019BAAS...51g.109H", "2019ICRC...36..836A", "2019JCAP...08..023F", "2019MNRAS.484.1790K", "2019MNRAS.484.2944G", "2019arXiv190305249R", "2019arXiv190701174S", "2020ApJ...905..178K", "2020IAUS..342..158T", "2020JPhCS1342a2104A", "2020MNRAS.491.4194L", "2020MNRAS.497.5318F", "2021ApJ...919..137T", "2021Galax...9...87L", "2021MNRAS.503.4032A", "2021PhRvD.104h3013A", "2022ApJ...934..158A", "2022ApJ...935..159K", "2022Galax..10...61R", "2022PhRvD.106j3021X", "2022icrc.confE.841U", "2023ApJ...944..157C", "2023JHEAp..40...55F", "2023MNRAS.524...76Z", "2023arXiv230211276Y" ]
[ "astronomy" ]
8
[ "galaxies: active", "galaxies: individual: NGC 1275", "M87", "Cen A", "neutrinos", "radiation mechanisms: non-thermal", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1968PhRvL..21.1016S", "1970ApJ...161L...1B", "1975ApJ...199L.139W", "1976ApJ...206L..45M", "1981ApJ...244..429B", "1984ARA&A..22..425H", "1986rpa..book.....R", "1991AJ....101.1632B", "1991ApJ...373L...1T", "1992ApJS...83...29S", "1994ApJ...430L..41V", "1994ApJ...430L..45W", "1994hea..book.....L", "1996A&A...309..375D", "1996ApJ...467..597B", "1996ApJ...473..254S", "1997ApJ...479..290S", "1997NIMPA.389...81B", "1997PhRvL..78.2292W", "1998A&A...330...97S", "1998ApJ...509..608T", "1998PhRvD..59b3002W", "1998tx19.confE..47W", "1999A&AS..139..545K", "1999APh....11..221S", "1999ApJS..123...79H", "1999Natur.401..891J", "2000ApJ...528..276M", "2000ApJ...530..233W", "2001APh....14..271S", "2001APh....15..121M", "2001ApJ...548..787S", "2001ApJ...561L..51P", "2001PhRvD..64b3002B", "2001PhRvL..87v1102A", "2002ApJ...564..683M", "2003ApJ...593..169H", "2004ApJ...609..539K", "2004PhRvD..70b3006B", "2005A&A...440..775K", "2005A&A...441..465A", "2005ApJ...622..910A", "2005ApJ...634L..33G", "2005NCimC..28..755K", "2005SSRv..120...95R", "2006ApJ...644..148P", "2006MNRAS.368.1500T", "2006PASJ...58..211H", "2006PASJ...58..261A", "2006PhRvD..74i4009B", "2006Sci...314.1424A", "2007AN....328..614U", "2007Ap&SS.309..407H", "2007ApJ...654..186F", "2007ApJ...655..781S", "2007ApJ...660..200L", "2007ApJ...665..209M", "2007ApJ...668L..27K", "2007ApJ...670L..81H", "2007Sci...318..938P", "2008A&A...478..111L", "2008A&A...487..837F", "2008APh....29..188P", "2008ApJ...689..775T", "2008PhR...458..173B", "2008PhRvD..78b3007C", "2009AJ....137.3718L", "2009APh....32...53H", "2009ApJ...690..367A", "2009ApJ...695L..40A", "2009ApJ...699...31A", "2009ApJ...706L.275A", "2009ApJ...707...55A", "2009JCAP...11..009L", "2009MNRAS.393.1041H", "2009NJPh...11f5016D", "2010ApJ...710..634A", "2010ApJ...710.1422R", "2010ApJ...719..459J", "2010ApJ...719.1433A", "2010MNRAS.405.2810H", "2010PhRvD..81l3001M", "2010PhRvD..82b2004G", "2010PhRvL.104i1101A", "2010arXiv1001.0059O", "2011MNRAS.413.2785B", "2011PhRvD..83i2003A", "2011PhRvD..84l2005A", "2011arXiv1107.4809T", "2012A&A...539L...2A", "2012ApJ...745..196R", "2012ApJ...746..141A", "2012ApJ...746..151A", "2012ApJ...749...63M", "2012ApJ...753...40F", "2012NIMPA.689...87A", "2012PhRvD..85d3012S", "2012PhRvD..85f2004K", "2013ApJ...766...73L", "2013ApJ...768L...1A", "2013ApJ...770L...6S", "2013PhRvD..88d7301S", "2013PhRvD..88j3003R", "2013PhRvL.111b1103A", "2013PhRvL.111l1102M", "2013Sci...342E...1I", "2014A&A...562A..12P", "2014ApJ...783...44F", "2014ApJ...785...53N", "2014ApJ...790L..21A", "2014JHEAp...3...29D", "2014MNRAS.437.2187F", "2014MNRAS.441.1209F", "2014MNRAS.445..570P", "2014PhRvD..90b3007M", "2014PhRvD..90b3010A", "2014PhRvL.113j1101A", "2014arXiv1410.8549M", "2015APh....70...54F", "2015APh....71....1F", "2015ATel.7856....1S", "2015ApJ...806..159K", "2015ApJ...815L..25G", "2015IAUS..313..175F", "2015JCAP...09..036T", "2015JETP..120..541K", "2015MNRAS.448.2412P", "2015MNRAS.451.1502T", "2015MNRAS.452.1877P", "2015arXiv150407592K", "2016ApJ...826...31F", "2016JHEAp...9...25F", "2016JHEAp..11...29F", "2016MNRAS.455..838K", "2016NatPh..12..807K" ]
[ "10.3847/0004-637X/830/2/81", "10.48550/arXiv.1607.04633" ]
1607
1607.04633_arXiv.txt
Radio galaxies (RGs) are radio loud active galaxy nuclei (AGN) exhibiting clear structure of a compact central source, large-scale jets and lobes. These sources are also of interest due to the close proximity to the earth affording us an excellent opportunity for studying the physics of relativistic outflows. They have been widely studied from radio wavelengths to MeV-GeV $\gamma$-rays and recently at very-high-energy (VHE) by Imagine Atmospheric Cherenkov Telescope (IACT). RGs are generally described by the standard non-thermal synchrotron self-Compton (SSC) model \citep{2010ApJ...719.1433A,2012ApJ...753...40F,1998ApJ...509..608T}. In SSC framework, low-energy emission; radio through optical, originates from synchrotron radiation while HE photons; X-rays through $\gamma$-rays, come from inverse Compton scattering emission. However, this model with only one population of electrons, predicts a spectral energy distribution (SED) that can be hardly extended up to higher energies than a few GeVs \citep{2005ApJ...634L..33G,2008A&A...478..111L}. In addition, some authors have suggested that the emission in the GeV - TeV energy range may have origins in different physical processes \citep{2011MNRAS.413.2785B,2005ApJ...634L..33G}.\\ The Telescope Array (TA) experiment, located in Millard Country (Utah), was designed to study ultra-high-energy cosmic rays (UHECRs) with energies above 57 EeV \citep{2012NIMPA.689...87A}. TA observatory, with a field of view covering the sky region above -10$^\circ$ of declination, detected a cluster of 72 UHECRs events centered at R. A.=146$^\circ$.7, Dec.=43$^\circ.2$. It had a Li-Ma statistical significance of 5.1$\sigma$ within 5 years of data taking \citep{2014arXiv1404.5890T}. The Pierre Auger observatory (PAO), located in the Mendoza Province of Argentina, was designed to determine the arrival directions and energies of UHECRs using four fluorescence telescope arrays and the High Resolution Fly's eye (HiRes) experiment, located in the west-central Utah desert, consisted of two detectors that observed cosmic ray showers via the fluorescence light. PAO studying the composition of the high-energy showers found that the distribution of their properties was located in somewhere between pure proton (p) and pure iron (Fe) at 57 EeV\citep{2008ICRC....4..335Y, 2008APh....29..188P, 2007AN....328..614U}, although the latest results favored a heavy nuclear composition \citep{2010PhRvL.104i1101A}. By contrast, HiRes data were consistent with a dominant proton composition at this energy range \citep{2007AN....328..614U, 2008ICRC....4..385E}. PAO detected 27 UHECRs, with two of them associated with Centaurus A (Cen A), being reconstructed inside a circle centered at the position of this FR I with a aperture of 3.1$^{\circ}$.\\ The IceCube detector located at the South Pole was designed to record the interactions of neutrinos. Encompassing a cubic kilometer of ice and almost four years of data taking (from 2010 to 2014), the IceCube telescope reported with the High-Energy Starting Events (HESE)\footnote{http://icecube.wisc.edu/science/data/HE-nu-2010-2014} catalog a sample of 54 extraterrestrial neutrino events in the TeV - PeV energy range. Arrival directions of these events are compatible with an isotropic distribution. The neutrino flux is compatible with a high component due to extragalactic origin \citep{2014PhRvD..90b3010A, 2015ApJ...815L..25G}. For instance, gamma-ray bursts \citep{2013PhRvL.111l1102M, 2013PhRvD..88j3003R, 2014MNRAS.437.2187F, 2013ApJ...766...73L, 2015JCAP...09..036T, 2014MNRAS.445..570P, 2016JHEAp...9...25F} and AGN (\cite{2014JHEAp...3...29D,2013PhRvD..88d7301S, 2015MNRAS.448.2412P, 2015MNRAS.452.1877P, 2014arXiv1410.8549M, 2015APh....71....1F, 2015APh....70...54F}), etc.\\ Hadronic processes producing VHE neutrinos and photons through the acceleration of cosmic rays in AGNs have been explored by many authors \citep{2001PhRvL..87v1102A, 2008PhR...458..173B, 2014arXiv1410.8124K, 2001APh....15..121M, 2008PhRvD..78b3007C, 2009NJPh...11f5016D, 2015MNRAS.451.1502T, 2015ApJ...806..159K, 2007ApJ...670L..81H, 2003ApJ...593..169H}. In particular, GeV-TeV $\gamma$-ray fluxes interpreted as pion decay products from photo-hadronic interactions occurring close to the core of RGs have been also discussed for different sources \citep{2012PhRvD..85d3012S, 2012ApJ...753...40F, 2015arXiv150407592K, 2016MNRAS.455..838K, 2014MNRAS.441.1209F, 2014A&A...562A..12P,2015IAUS..313..175F}.\\ In this work we introduce a leptonic and hadronic model to describe the broadband SED of the closest RGs. For the leptonic model, we present the SSC model to explain the SED up to dozens of GeV and for the hadronic model, we propose that the proton-photon (p$\gamma$) interactions occurring close to the core could describe the $\gamma$-ray spectrum at the GeV - TeV energy range. Correlating the TeV $\gamma$-ray, UHECR and neutrino spectra through p$\gamma$ interactions, we estimate their fluxes and number of events expected in PAO, TA and HiRes experiments and IceCube telescope, respectively. For this correlation, we have assumed that the proton spectrum is extended through a simple power law up to UHEs.
We have proposed a leptonic and hadronic model to explain the broadband SED spectrum observed in the closest RGs. In the leptonic model, we have used the SSC emission to describe the SED up to dozens of GeV. To explain the spectrum from hundreds of GeV up to a few TeV, we have introduced the hadronic model assuming that accelerated protons in the inner jet interact with the photon population at the SED peaks. Evoking these interaction, we have interpreted the TeV $\gamma$-ray spectra as $\pi^0$ decay products in Cen A, M87 and NGC 1275.\\ Correlating the TeV $\gamma$-ray and HE neutrino fluxes through p$\gamma$ interactions, we have computed the HE and UHE neutrino fluxes, and the neutrino event rate expected in a kilometric scale neutrino detector. The neutrino event rate was obtained through MC simulation by considering a region of 1$^{\circ}$ around the source position and assuming a hypothetical Km$^3$ Cherenkov telescope. We found that the neutrino fluxes produced by p$\gamma$ interactions close to the core of RGs cannot explain the astrophysical flux and the expected $\nu_{\mu}$ events in a neutrino telescope are consistent with the nonneutrino track-like associated with the location of the closest RGs \citep{2013PhRvL.111b1103A, 2013Sci...342E...1I, 2014PhRvL.113j1101A, 2015ATel.7856....1S}. In addition, the atmospheric muon neutrino background is also shown.\\ Extrapolating the proton spectrum by a simple power law up to UHEs, we have computed the number of UHECRs expected from Cen A, M87 and NGC1275. We found that those UHECRs obtained with our model is in agreement with the TA and PAO observations. Although UHECRs from Cen A can hardly be accelerated up to the PAO energy range at the emitting region (E$_{max}$= 40 EeV), they could be accelerated during the flaring intervals, with small changes in the strength of magnetic field and/or emitting region and in the giant lobes. It is very interesting the idea that UHECRs could be accelerated partially in the jet at energies ($<40\times 10^{19}$ eV) and partially in the Lobes at ($E>40\times 10^{19}$ eV) \citep{2014ApJ...783...44F}. If UHECRs have a heavy composition as suggested by PAO \citep{2010PhRvL.104i1101A}, UHE heavy nuclei in Cen A could be accelerated at energies greater than $\sim$ 80 EeV, thus reproducing the detections reported by PAO. It is worth noting that if radio Galaxies are the sources of UHECRs, their on-axis counterparts (i.e. blazars, and flat-spectrum radio quasars) should be considered a more powerful neutrino emitters \citep{2001PhRvL..87v1102A, 2014JHEAp...3...29D,2014PhRvD..90b3007M}. In fact, a PeV neutrino shower-like was recently associated with the flaring activity of the blazar PKS B1424-418 \citep{2016arXiv160202012K}. \\ In summary, we have showed that leptonic SSC and hadronic processes are required to explain the $\gamma$-ray fluxes at GeV- TeV energy range. We have successfully described the TeV $\gamma$-ray \citep{2009ApJ...695L..40A, 2006Sci...314.1424A, 2010ApJ...710..634A}, HE neutrinos \citep{2013PhRvL.111b1103A, 2013Sci...342E...1I, 2014PhRvL.113j1101A, 2015ATel.7856....1S} and UHECRs \citep{2011arXiv1107.4809T, 2013ApJ...768L...1A} around the closest RGs.
16
7
1607.04633
The closest radio galaxies; Centaurus A (Cen A), M87, and NGC 1275, have been detected from radio wavelengths to TeV γ-rays, and also studied as high-energy neutrino and ultra-high-energy cosmic-ray (UHECR) potential emitters. Their spectral energy distributions (SEDs) show a double-peak feature, which is explained by a synchrotron self-Compton (SSC) model. However, TeV γ-ray measured spectra could suggest that very-high-energy γ-rays might have a hadronic origin. We introduce a lepto-hadronic model to describe the broadband SED; from radio to sub-GeV photons as synchrotron SSC emission and TeV γ-ray photons as neutral pion decay resulting from pγ interactions occurring close to the core. These photo-hadronic interactions take place when Fermi-accelerated protons interact with the seed photons around synchrotron SSC peaks. Obtaining a good description of the TeV γ-ray fluxes, first, we compute neutrino fluxes and events expected in the IceCube detector and, second, we estimate UHECR fluxes and the event rate expected in Telescope Array, Pierre Auger, and HiRes observatories. Within this scenario, we show that the expected high-energy neutrinos cannot explain the astrophysical flux observed by IceCube, and the connection with UHECRs observed by Auger experiment around Cen A might be possible only considering a heavy nuclei composition in the observed events.
false
[ "synchrotron SSC peaks", "TeV γ-ray", "synchrotron SSC emission", "TeV γ-ray measured spectra", "neutrino fluxes", "sub-GeV photons", "UHECR fluxes", "Cen A", "rays", "Auger experiment", "SSC", "Pierre Auger", "radio wavelengths", "HiRes observatories", "events", "Centaurus A", "TeV γ-ray photons", "TeV", "neutral pion decay", "pγ interactions" ]
5.710986
0.425503
13
12620348
[ "Dietrich, Tim", "Ujevic, Maximiliano", "Tichy, Wolfgang", "Bernuzzi, Sebastiano", "Brügmann, Bernd" ]
2017PhRvD..95b4029D
[ "Gravitational waves and mass ejecta from binary neutron star mergers: Effect of the mass ratio" ]
148
[ "Max Planck Institute for Gravitational Physics, Albert Einstein Institute, D-14476 Golm, Germany", "Centro de Ciências Naturais e Humanas, Universidade Federal do ABC, 09210-170 Santo André, São Paulo, Brazil", "Department of Physics, Florida Atlantic University, Boca Raton, Florida 33431, USA", "DiFeST, University of Parma, and INFN Parma I-43124 Parma, Italy", "Theoretical Physics Institute, University of Jena, 07743 Jena, Germany" ]
[ "2016CQGra..33x4004E", "2017ApJ...842L..10R", "2017ApJ...846...62S", "2017ApJ...850L..24G", "2017ApJ...850L..39A", "2017ApJ...851L..16A", "2017CQGra..34h4002P", "2017CQGra..34j5014D", "2017CQGra..34x5003C", "2017PhRvD..95d4045D", "2017PhRvD..95l4006D", "2017PhRvD..96d3019K", "2017PhRvD..96f3011M", "2017PhRvD..96l4005B", "2017PhRvD..96l4035C", "2017PhRvL.119p1101A", "2018ApJ...852L..29R", "2018ApJ...855...67L", "2018ApJ...856..101M", "2018ApJ...858...74M", "2018ApJ...861...55M", "2018ApJ...865L..21K", "2018ApJ...866...60P", "2018ApJ...867...39K", "2018ApJ...867...95H", "2018ApJ...869..130R", "2018CQGra..35xLT01D", "2018Galax...6..119N", "2018IJMPD..2742005H", "2018JApA...39...45H", "2018MNRAS.481.3670R", "2018PhRvD..97f4002D", "2018PhRvD..98j4005C", "2018PhRvL.120k1101Z", "2018PhRvL.120v1101D", "2018arXiv180706857V", "2019ARNPS..69...41S", "2019AnPhy.41167958B", "2019ApJ...876..139G", "2019ApJ...877....2W", "2019ApJ...881...89L", "2019ApJ...884...40M", "2019ApJ...887L..35B", "2019ChPhC..43e4108C", "2019EPJA...55..203S", "2019JPhG...46k3002B", "2019LRR....23....1M", "2019Parti...2...44H", "2019Parti...2..365G", "2019PhRvC.100c5803A", "2019PhRvD..99b4029D", "2019PhRvD..99d4051A", "2019PhRvD.100d3001P", "2019PhRvD.100j4029B", "2019PhRvD.100l4046T", "2019PhRvL.123d1102Z", "2019PrPNP.10903714B", "2019Univ....5..156H", "2020ARNPS..70...95R", "2020ApJ...889..171K", "2020ApJ...892...35N", "2020ApJ...897..150D", "2020ApJ...902L..12Z", "2020CQGra..37o5005W", "2020GReGr..52..108B", "2020IJMPD..2941015F", "2020JPhCS1668a2029N", "2020MNRAS.495.3780N", "2020MNRAS.497..643D", "2020MNRAS.497.1488B", "2020PhR...886....1N", "2020PhRvD.101f4052D", "2020PhRvD.101h4039V", "2020PhRvD.101j3008C", "2020PhRvD.101j3036G", "2020PhRvD.102b4087C", "2020PhRvD.102d4040X", "2020PhRvL.125z1101H", "2020Sci...370.1450D", "2020Symm...12.1249R", "2020arXiv200509618H", "2021A&G....62.4.15N", "2021ApJ...908..122G", "2021ApJ...908..194T", "2021ApJ...911...97M", "2021ApJ...912...80M", "2021ApJ...913..100K", "2021ApJ...915..108N", "2021ApJ...916L..13I", "2021ApJ...918...10W", "2021ApJ...920L...3C", "2021ApJ...921..161D", "2021ApJ...922L..19T", "2021EPJST.230..543H", "2021GReGr..53...27D", "2021JPlPh..87a8402A", "2021LRR....24....5K", "2021MNRAS.500.1817L", "2021MNRAS.502.1843N", "2021MNRAS.505.3016N", "2021MNRAS.508.1732K", "2021PhRvD.103l4015G", "2021PhRvD.104h3004Z", "2021PhRvD.104h3029P", "2022ApJ...932L...7C", "2022ApJ...933...22K", "2022ApJ...937...79C", "2022JCAP...02..027O", "2022MNRAS.510.2968K", "2022MNRAS.513.3646P", "2022MNRAS.516.4760C", "2022NatAs...6..961A", "2022Parti...5..198B", "2022PhRvD.105f4050D", "2022PhRvD.105j4019W", "2022PhRvD.106b3029U", "2022PhRvD.106b4001D", "2022PhRvD.106d4026K", "2022PhRvD.106h4039D", "2022PhRvD.106j3020T", "2022PhRvD.106j3027K", "2022PhRvL.129c2701P", "2022Univ....8..633W", "2022arXiv221207498J", "2023ApJ...958L..33G", "2023CQGra..40h5011G", "2023MNRAS.520.2727N", "2023MNRAS.522.5848L", "2023MNRAS.525.3384K", "2023NewA..10402067S", "2023PhRvD.107b4025U", "2023PhRvD.107f3028H", "2023PhRvD.107j3053M", "2023arXiv230600948B", "2024JPhCS2742a2009R", "2024MNRAS.527.8043C", "2024MNRAS.527.8812J", "2024MNRAS.528.4785Z", "2024MNRAS.528.5836R", "2024PhRvD.109d3015B", "2024PhRvD.109f4009B", "2024PhRvD.109l3011S", "2024Sci...383..275B", "2024arXiv240103750B", "2024arXiv240403156Z", "2024arXiv240509513S", "2024arXiv240513687K", "2024arXiv240605211T" ]
[ "astronomy", "physics" ]
19
[ "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.1103/PhysRevD.95.024029", "10.48550/arXiv.1607.06636" ]
1607
1607.06636_arXiv.txt
\label{sec:intro} Binary neutron star (BNS) mergers are primary sources of gravitational waves (GWs) and are associated with a variety of electromagnetic (EM) emissions. GW observations of BNS are eagerly expected in the upcoming LIGO-Virgo runs, after the first binary black hole (BBH) GW detections GW150914~\cite{Abbott:2016blz} and GW151226~\cite{Abbott:2016nmj}. Such GWs will allow us to place constraints on the nature of matter at densities above nuclear density, e.g.~\cite{DelPozzo:2013ala}, and to identify the origin of EM emissions like kilo/macronovae and short gamma-ray burst (SGRB), e.g.~\cite{Rosswog:2015nja,Fernandez:2015use}. Kilo/macronovae events are transient emissions in the optical or near-infrared band observed in e.g.~\cite{Tanvir:2013pia}. They are believed to be triggered by the radioactive decay of r-process nuclei in the neutron-rich material ejected during a BNS merger. SGRBs models are instead based on highly relativistic outflows, powered e.g.~by the merger remnant accretion disk~\cite{Paczynski:1986px,Eichler:1989ve}. Their joint observation with GWs might be challenging due to the short duration of the burst and to the fact that they are highly collimated emissions~\cite{Soderberg:2006bn}. Additionally, the interaction of mildly or sub- relativistic outflows with the surrounding material generates synchrotron radiation known as radio flares~\cite{Nakar:2011cw}. This emission can persist from months to years after merger, depending on the composition, which makes radio flares a particularly attractive EM counterpart to detect. Understanding the dependency of the GW and EM emissions on the source's parameters is of key importance for GW astronomy and multimessenger astrophysics. The BNS parameter space is composed by the component masses (and spins), and by a choice of equation of state (EOS) describing the NS matter. The ranges of the mass and spins parameters are rather uncertain. The expected NS mass range is $\sim 0.9 - 3 M_\odot$. The lower bound is inferred from the formation scenario (gravitational-collapse) and from current observations, although those measurements have typically large uncertainties, see e.g.~\cite{Rawls:2011jw,Ozel:2012ax}. The upper bound is inferred from stability argument (maximum theoretical mass), from precise measurements of $\sim 2M_\odot$ NSs in double NS systems~\cite{Demorest:2010bx,Antoniadis:2013pzd}, and from models of SGRBs, which suggest a maximum mass of $2.2M_\odot$, e.g.~\cite{Lawrence:2015oka}. There exist observations of larger NS masses but with large uncertainties, so they cannot give strong constraints on the NS maximum mass \cite{Lattimer:2012nd}. Also, the maximum mass of a NS is determined by the particular EOS. Most tabulated EOS compatible with astrophysical constraints support NSs with maximum masses in a range of $\sim 2-3M_\odot$. The above considerations suggest that the BNS mass-ratio \be q := M^A/M^B\geq1 \ , \ee where $M^{A,B}$ are the individual gravitational masses of the NSs (in isolation), is most likely constrained to $q\lesssim 2-3$. Observations suggest that BNS systems consist of equal mass NSs with masses of around $\sim 1.35 M_\odot$ and that the mass-ratio is close to one e.g.~\cite{Swiggum:2015yras,Lattimer:2012nd}. However, only approximately a dozen BNS systems are known so far and only $6$ of these systems have well determined masses and will merge within a Hubble time. The lack of very unequal mass configurations might only be a selection effect. For example Ref.~\cite{Martinez:2015mya} discovered a compact binary system with a mass ratio of $q \approx 1.3$, suggesting that BNS with larger mass ratios exist. Population synthesis models for binaries formed ``\emph{in situ}'' predict a wider range of masses and mass ratios up to $q\approx 1.9$~\cite{Dominik:2012kk,Dietrich:2015pxa}. Also NS spins are constrained by theoretical arguments and observations, e.g. \cite{Burgay:2003jj,Lynch:2011aa,Miller:2014aaa,Dietrich:2015pxa}. However, we do not consider here the NS rotation and we remind to the above references and to future work \cite{Dietrich:2016prep2}. Parameter space investigation of BNS are challenging due to the unknown EOS, and the need of simulating each masses (and spins) configuration with different EOS. Most numerical relativity studies of BNS systems have focused on equal masses and irrotational configurations. The first simulations of unequal mass systems have been presented in \cite{Shibata:2003ga,Rezzolla:2010fd,Hotokezaka:2012ze} using polytropic and piecewise polytropic EOS. Mass ejecta in $q\neq1$ BNS simulations have been studied in e.g.~\cite{Bauswein:2013yna,Hotokezaka:2012ze,Dietrich:2015iva}. Unequal mass simulations with microphysical EOS and neutrino cooling have been presented in e.g.~\cite{Bernuzzi:2015opx,Lehner:2016wjg,Lehner:2016lxy} and with radiation-hydrodynamics in~\cite{Sekiguchi:2016bjd}. Previous works were restricted to mass ratios $q\leq 1.35$. Overall, the main results are that i) asymmetric mergers produce more massive ejecta with smaller electron fractions than the corresponding equal mass setups ii) unequal mass systems are likely to produce kilonovae and iii) the remnant disk mass increases for an increasing mass ratio. In \cite{Dietrich:2015pxa} we reported an upgrade of the SGRID code able to generate generic initial data for BNS simulations together with few preliminary evolutions. Among other results, we showed the possibility of generating ``large mass-ratio'' configurations with $q\sim2$. The test evolution of a $q=2$ BNS showed interesting features, including large mass ejection and mass transfer from one star to the other during the last revolutions. In this work we study the effect of the binary's mass-ratio $q$ on the GWs and on the characteristics of possible EM emission associated to dynamical mass ejecta, in particular macronovae and radio flares. We present a new set of (3+1)D numerical relativity simulations of the merger and postmerger phase, and focus on a previously inaccessible region of the binary parameter space spanning $q\in[1,1.75]$ for different masses and equations of state, and a setup with $q=2$. The article is structured as follows: In Sec.~\ref{sec:simeth}, we give a short description of the numerical methods and describe important quantities used to analyze our simulations. Section~\ref{sec:config} summarizes our configurations and the investigated part of the BNS parameter space. Section~\ref{sec:dynamics} deals with the dynamics of the simulation, where in particular we focus on the mass-transfer, the ejecta, and the final remnant. The GW signal is investigated in Sec.~\ref{sec:GW} with respect to spectrograms, the sky location and the emitted GW energy per mode. Sec.~\ref{sec:EM} focuses on EM counterparts and we conclude in Sec.~\ref{sec:conclusion}. In Appendix~\ref{sec:accuracy} we test the accuracy of our simulations with respect to conserved quantities, convergence, the constraints. Throughout this work we use geometric units, setting $c=G=M_\odot=1$, though we will sometimes include $M_\odot$ explicitly or quote values in CGS units for better understanding. Spatial indices are denoted by Latin letters running from 1 to 3 and Greek letters are used for spacetime indices running from 0 to 3.
\label{sec:conclusion} In this article we studied the effect of the mass-ratio by a large set of new numerical relativity simulations spanning, for the first time, up to values $q\sim2$. Our findings are summarized in what follows. \\ \paragraph*{Mass transfer:} A resolution study of simulations with $q=2$ showed that mass transfer during the last orbits is very dependent on the grid resolution. In particular, by increasing resolution the amount of transfered mass decreases. We conclude that no significant mass transfer happens during the merger of BNS, even in cases with large mass-ratio. \\ \paragraph*{Mass ejection:} Mass ejection in large $q$ systems is primary due to a centrifugal effect and originates from the companion's tidal tail (or partial disruption) during the late inspiral and merger. Ejecta components due to shock-driven ejecta are only dominant for configurations with a soft EOS and rather independent of the mass ratio. We showed, for the first time in the context of BNS, that the dependence of the ejecta mass and kinetic energy is essentially linear on the mass-ratio $q$ for stiff EOSs, see Fig.~\ref{fig:M_ej(q)}. Also the velocity of the ejecta depend significantly on the mass ratio. In particular, the component perpendicular to the orbital plane decreases for increasing $q$, since torque becomes the dominant ejecta mechanism. For large $q$ the ejecta are almost entirely about the orbital plane. Overall, these ejecta properties for large $q$ resamble those of black hole - neutron star binaries and lead to characteristic features of electromagnetic counterparts (see below). Finally, the total mass of the configuration plays a minor role and is less important than the EOS or the mass-ratio. \\ \paragraph*{Merger remnant:} The lifetime of the merger remnant depends strongly on the EOS, in most cases softer EOSs lead to an earlier collapse. The mass-ratio is a secondary effect, but larger $q$ lead to delay collapse. We also showed that in most cases for which a black hole forms the rest mass of the accretion disk increases. In cases of a prompt collapse no massive accretion disk forms. \\ \paragraph*{Gravitational Waves:} Varying the mass ratio leads to quantitative changes to the GW frequency, and to qualitative changes of the postmerger spectra. The GW merger frequency is in general largest for equal masses and decreases for increasing $q$. For MS1b $f_{\rm mrg}$ decreases from 1.45~kHz to 0.09~kHz when $q$ goes from $1$ to $1.75$ respectively. No significant effect are instead on the postmerger frequency $f_2$. We believe the latter is due to the limited accuracy the peak frequency can be extracted, ultimately related to the broad character of the spectra peak. However, we find that for unequal masses the characteristic secondary peaks $f_s$ in the spectrum tend to disappear for large $q$, and they are actually absent for high mass ratio systems, see Fig.~\ref{fig:hspectra}. Our spectrograms highlight the rich structure of the multipolar GW waveform. Modes with azimuthal number $m=1,3$, in particular, become progressively more relevant for larger $q$, although it is unlikely their effect will be observed in next LIGO/Virgo observations, e.g. \cite{Radice:2016gym}. Furthermore, for configurations producing a MNS merger remnant, the spectrograms show that the postmerger frequencies increase over time in a chirp-like fashion as the merger remnant becomes more compact. This implies that the postmerger spectrum is in fact continuous and not discrete as anticipated in \cite{Bernuzzi:2015rla}. The total emitted GW energy is a decreasing function of $q$ during both the inspiral and the post-merger phase. This qualitative behavior is already known from binary black hole systems, but here we extend the result to BNS. This is nontrivial since for BNS tidal interactions play an important role during inspiral-merger and the postmerger phase has different physics from black hole binaries. We find that neutron stars with softer EOS emit more energy. In addition to the total energy, also the mode hierarchy changes by varying $q$. We find that the energy in the $m=3$ (and $m=1$) emission channel increase for larger $q$ and contribute up to $1\%$ to the total emitted energy up to merger. \\ \paragraph*{Electromagnetic counterparts:} A GW detection of BNSs will trigger observations to capture follow-up electromagnetic emissions. We used simplified models to estimate the luminosity, peak time, and light curves of macronovae counterparts and the peak time and fluence of radio flares. We showed that the peak luminosity, peak time, and persistency of these counterparts are strongly dependent on the mass ratio $q$. Unequal mass BNS systems are more luminous in the EM, than equal mass systems because of more massive ejecta. The larger the mass ratio, the more delayed is the luminosity peak. Also our estimated macronova lightcurves are more persistent for larger mass ratios; they could be detected up to a few weeks, see Fig. \ref{fig:lightcurve}. Similarly to black hole neutron star mergers and differently from $q\sim1$ BNS, the dynamical ejecta of large $q$ BNS are confined to the equatorial plane and will not obscure optical emissions from the disk wind~\cite{Kasen:2014toa}. Thus, the latter might be detectable for face-on binaries \footnote{We thank D.Radice for pointing this out.}. \\ Numerical uncertainties have been carefully evaluated on multiple quantities, see Appendix~\ref{sec:accuracy}, and we are confident on the presented results. However, our simulations do not account for microphysics as done in other works e.g.~\cite{Just:2014fka,Bernuzzi:2015opx,Palenzuela:2015dqa,Lehner:2016wjg,Lehner:2016lxy,Sekiguchi:2016bjd,Radice:2016dwd}. The simplified assumptions in our work have been necessary to simulated a large number of BNS configurations, and in order to better control the numerical errors. We believe our results hold at least at a qualitative level; dedicated simulations of selected configurations that include a more sophisticated treatment of microphysics should be performed in the future to validate our predictions. \appendix
16
7
1607.06636
We present new (3 +1 )D numerical relativity simulations of the binary neutron star (BNS) merger and postmerger phase. We focus on a previously inaccessible region of the binary parameter space spanning the binary's mass ratio q ∼1.00 - 1.75 for different total masses and equations of state, and up to q ∼2 for a stiff BNS system. We study the mass ratio effect on the gravitational waves (GWs) and on the possible electromagnetic (EM) emission associated with dynamical mass ejecta. We compute waveforms, spectra, and spectrograms of the GW strain including all the multipoles up to l =4 . The mass ratio has a specific imprint on the GW multipoles in the late-inspiral-merger signal, and it affects qualitatively the spectra of the merger remnant. The multipole effect is also studied by considering the dependency of the GW spectrograms on the source's sky location. Unequal mass BNSs produce more ejecta than equal mass systems with ejecta masses and kinetic energies depending almost linearly on q . We estimate luminosity peaks and light curves of macronova events associated with the mergers using a simple approach. For q ∼2 the luminosity peak is delayed for several days and can be up to 4 times larger than for the q =1 cases. The macronova emission associated with the q ∼2 BNS is more persistent in time and could be observed for weeks instead of a few days (q =1 ) in the near infrared. Finally, we estimate the flux of possible radio flares produced by the interaction of relativistic outflows with the surrounding medium. Also in this case a large q can significantly enhance the emission and delay the peak luminosity. Overall, our results indicate that the BNS merger with a large mass ratio has EM signatures distinct from the equal mass case and more similar to black hole-neutron star binaries.
false
[ "q", "equal mass systems", "dynamical mass ejecta", "different total masses", "BNS", "ejecta masses", "luminosity peaks", "Unequal mass BNSs", "several days", "postmerger phase", "kinetic energies", "a large q", "the q ∼2", "time", "a large mass ratio", "∼2", "EM signatures", "black hole-neutron star binaries", "the BNS merger", "the equal mass case" ]
6.558908
2.434398
62
12406371
[ "Vibert, D.", "Peillon, C.", "Lamy, P.", "Frazin, R. A.", "Wojak, J." ]
2016A&C....17..144V
[ "Time-dependent tomographic reconstruction of the solar corona" ]
20
[ "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille), UMR 7326, 13388, Marseille, France", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille), UMR 7326, 13388, Marseille, France", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille), UMR 7326, 13388, Marseille, France", "Department of Atmospheric, Oceanic and Space Sciences, University of Michigan, Ann Arbor, MI 48109, USA", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille), UMR 7326, 13388, Marseille, France" ]
[ "2017SoPh..292...60L", "2017SoPh..292...97W", "2017SoPh..292..153L", "2017arXiv170605116W", "2019ApJ...887...83S", "2019ApJS..242....3M", "2019SSRv..215...39L", "2019SoPh..294...81V", "2019SoPh..294..162L", "2020A&A...642A...2R", "2020A&A...642A..13H", "2020ApJ...893...57M", "2020BAAA...61C..34V", "2020SoPh..295...89L", "2021ApJ...922..165M", "2022ApJ...928..131J", "2022JGRA..12730406L", "2022JSWSC..12...30B", "2022SoPh..297..120V", "2023SoPh..298..135A" ]
[ "astronomy" ]
10
[ "Sun: corona", "Sun: coronal structures", "Tomography", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1016/j.ascom.2016.09.001", "10.48550/arXiv.1607.06308" ]
1607
1607.06308_arXiv.txt
Progress in understanding the physics of the corona and the solar wind depends upon empirical constraints on the three-dimensional (3D) distributions of the plasma properties such as temperature, density and magnetic field in the Sun's corona. Solar rotational tomography (SRT) uses coronal images observed at multiple aspect angles thanks to solar rotation to provide 3D reconstructions of the corona assuming that it does not vary in time, at least during the time interval necessary to achieve a complete view (typically half a solar rotation). When the images are at soft X-ray or EUV wavelengths, one can reconstruct both the temperature and the density in 3D from about \SIrange{1.03}{1.25}{\Rsun} (the lower bound is greater than \SI{1.0}{\Rsun} due to optical depth effects encountered in some spectral lines), via a process called differential emission measure tomography (DEMT), see \citet{Frazin09} and \citet{Barbey13}. When the input images are coronagraphic white-light images of the K-corona, one may reconstruct the density at greater heights. For example, when the image data are from the LASCO-C2 coronagraph, the reconstruction region is roughly from \SIrange{2.4}{5.5}{\Rsun} \citep{Frazin02} and when the source is STEREO/COR1 the reconstruction region is between about \SIrange{1.5}{4.0}{\Rsun} \citep{Kramar09,Kramar14}. In addition, the use of white-light coronagraph data has the potential of taking advantage of both the polarized brightness ($pB$) and total brightness ($B$) as independent sources of information \citep{Frazin10}. Whereas the determination of the 3D distribution of the coronal density (as well as temperature) is a highly desirable objective, SRT has fundamental limitations that have prevented it from becoming broadly accepted as a useful tool for solar science. Calibration uncertainty in the image data is important \citep{Frazin12}, the finite \fov{} of the imaging instrument causes ``pile-up'' artifacts near the upper boundary \citep{Frazin02,Frazin10}, and the $\approx \SI{7}{\degree}$ tilt of the Sun's rotational pole relative to the ecliptic is theoretically a source of non-uniqueness, but the dominant source of error and uncertainty is the inherently dynamic nature of the solar corona so that it evolves as the Sun rotates. Thus, when one sees a movie of the corona, it is difficult to disentangle effects of the coronal dynamics from those of solar rotation, leading to a fundamental ambiguity and very serious artifacts in tomographic reconstructions of the corona as first shown by \citet{Frazin05a}. Since then, a new method to perform time-dependent tomography involving Kalman filters was developed by \mbox{\citet{Butala10}} and applied to STEREO/COR1 images. It did improve the quality of the reconstruction by reducing the number of artifacts; however, and as we shall discuss later on, it faced several difficulties leaving the time-dependent tomography problem un\-der-de\-ter\-mi\-ned and the solution very reliant on the regularization choices. Moreover the complexity of the Kalman filter turns out to be of marginal utility due to the fact that we only have uninformative dynamical models of the corona at our disposal, whereas a time-dependent tomography can be achieved by a simpler multiple regularization approach as proposed in this article. Interestingly, the EUV tomography problem appears to be less compromised by this issue, most likely due to the fact that lines-of-sight that hit the solar disk (as well as those above the limb, as in the case of coronagraph data) are included in the inversion, which tend to stabilize the solution. Thus, to date, most scientific results from SRT have been produced with EUV tomography \citep{Vasquez09,Vasquez10,Vasquez11,Huang12,Nuevo13}. This article introduces several new strategies intended to mitigate the undesirable effects of coronal dynamics on tomographic reconstructions based on white-light coronagraph images. These new methods involve two different approaches: \begin{enumerate} \item For the part of the reconstruction located at a given radius $r_0$ and above, the goal is to de-emphasize the influence of coronal dynamics taking place at inferior radii $(r < r_0)$. The rationale behind this is that, at large heights, the corona is less dense and tends to be less dynamic. Thus, it is hoped that the reconstruction at $r > r_0$ will not be ``contaminated'' by the contribution of material with \los{} at $r < r_0$. These procedures involve weighting and masking schemes. Part of the justification for this approach is that, strictly speaking, one does not need projections inferior to $r_0$ to reconstruct the part of the object at $r_0$ and above \citep{Frazin10,Louis83}. \item In using Kalman filters or \spatiotemp{} regularization for tomography of the solar corona, it has been observed that, unless very strong temporal regularization is used (greatly reducing the temporal variability of the solution), the reconstructed density tends to be concentrated in the plane that contains the Sun's center and is perpendicular to the \los{} of the observation that was made at time $t$. Thus, the solution tracks the observation angle as the Sun rotates. The new strategy introduced to mitigate this effect is a novel type of \spatiotemp{} regularization, not based on gradients (as are most regularization operators), but instead specifically designed to suppress this rotational mode in the solution. \end{enumerate} The lessons learned from these new strategies are likely applicable to tomography based on EUV images, but that is not explored here. The article is organized as follows: first, a general explanation of the tomographic reconstruction of the solar corona is given in Section~\ref{theory}. Then improved static reconstructions are described in Section~\ref{static}. Next, time-dependent tomographic reconstruction methods are described in Section~\ref{dynamic} and applied to model and real images. Finally, we conclude in Section~\ref{conclusion}. \begin{tolerant}{6000}
\label{conclusion} In this article, we considered the main problem crippling solar rotational tomography, namely the assumption of a stable corona whereas it is intrinsically dynamic even during the minima of solar activity. Using a dynamic MHD model of the corona, we showed that a static reconstruction of the three-dimensional electron density results in severe artifacts that do not appear when applied to a static corona. We addressed the problem using two different approaches: i) mitigation of the temporal variation in the framework of a regularized static inversion, and ii) a true time-dependent inversion. Our main results based on tests performed on synthetic images constructed from this dynamic MHD model and on real LASCO-C2 images of the corona are summarized below. \begin{itemize} \item Crucial to testing our procedure and properly tuning the regularization parameters was the introduction of a time-dependent MHD model of the corona based on observed magnetograms to build a time-series of synthetic images of the corona. \item Our mitigation procedures --- multiple masking with simple juxtaposition and recursive combination of solutions, radial weighting with a mean radial intensity profile or an exponential profile --- do not convincingly improve the reconstruction. \item A true time-dependent inversion with \spatiotemp{} regularization does improve the situation and convincingly reduces the artifacts by a factor 30 and the normalized RMS error by a factor 2. \item The introduction of an additional \spatiotemp{} regularization that penalizes the coronal structures that are apparently fixed for the observer turns out to be disappointing. However it could potentially improve the reconstruction when using more frequent images. \item Whereas non-physical densities, although greatly reduced, are still present, our dynamic reconstruction appears qualitatively superior, exhibiting a generally smoother and more connected, \ie more physically reasonable, reconstructed streamer belt. \end{itemize} Future work will consider a possible improvement of the co-rotation regularization and will ultimately consist in applying our time-dependent SRT procedure to the whole set of LASCO-C2 images presently extending over 20 years, that is almost two solar cycles. It is hoped that the resulting four-dimensional estimates of the electron density will provide insight into the coronal and solar wind processes.
16
7
1607.06308
Solar rotational tomography (SRT) applied to white-light coronal images observed at multiple aspect angles has been the preferred approach for determining the three-dimensional (3D) electron density structure of the solar corona. However, it is seriously hampered by the restrictive assumption that the corona is time-invariant which introduces significant errors in the reconstruction. We first explore several methods to mitigate the temporal variation of the corona by decoupling the "fast-varying" inner corona from the "slow-moving" outer corona using multiple masking (either by juxtaposition or recursive combination) and radial weighting. Weighting with a radial exponential profile provides some improvement over a classical reconstruction but only beyond ≈ 3R<SUB>⊙</SUB>. We next consider a full time-dependent tomographic reconstruction involving spatio-temporal regularization and further introduce a co-rotating regularization aimed at preventing concentration of reconstructed density in the plane of the sky. Crucial to testing our procedure and properly tuning the regularization parameters is the introduction of a time-dependent MHD model of the corona based on observed magnetograms to build a time-series of synthetic images of the corona. Our procedure, which successfully reproduces the time-varying model corona, is finally applied to a set of 53 LASCO-C2 pB images roughly evenly spaced in time from 15 to 29 March 2009. Our procedure paves the way to a time-dependent tomographic reconstruction of the coronal electron density to the whole set of LASCO-C2 images presently spanning 20 years.
false
[ "inner corona", "synthetic images", "time", "reconstructed density", "multiple aspect angles", "multiple masking", "observed magnetograms", "recursive combination", "Solar rotational tomography", "LASCO-C2 images", "the solar corona", "white-light coronal images", "the time-varying model corona", "significant errors", "the coronal electron density", "a full time-dependent tomographic reconstruction", "a time-dependent tomographic reconstruction", "the corona", "a co-rotating regularization", "≈ 3R<SUB>⊙</SUB" ]
12.661941
15.461793
2
1492327
[ "Shannon, Andrew", "Bonsor, Amy", "Kral, Quentin", "Matthews, Elisabeth" ]
2016MNRAS.462L.116S
[ "The unseen planets of double belt debris disc systems" ]
36
[ "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Astrophysics Group, University of Exeter, Physics Building, Stocker Road, Exeter EX4 4QL, UK" ]
[ "2017MNRAS.465.2595M", "2017MNRAS.466.3973M", "2018A&A...611A..43L", "2018A&A...617A..76C", "2018ARA&A..56..541H", "2018MNRAS.475.4953R", "2018MNRAS.479.3300K", "2018MNRAS.479.5423M", "2018MNRAS.480.2757M", "2018MNRAS.481.5180M", "2018arXiv180105850C", "2018exha.book.....P", "2018haex.bookE.165K", "2019A&A...625A..21B", "2019MNRAS.484.1257M", "2019MNRAS.485.5511S", "2020A&A...635A..19E", "2020A&A...639A..54L", "2020A&A...642A..18L", "2020ApJ...898..146M", "2021A&A...645A..88M", "2021AJ....161...78M", "2021ApJ...910...13S", "2021MNRAS.500.1604B", "2021MNRAS.502.5390P", "2021MNRAS.503.1276M", "2021MNRAS.506.1978L", "2021exbi.book...15W", "2022A&A...659A.135P", "2022A&A...664A.139D", "2022MNRAS.512.4441F", "2022MNRAS.517.2546L", "2022MNRAS.517.5835S", "2023A&A...679A..58M", "2023ApJ...954..100S", "2023MNRAS.525L..36L" ]
[ "astronomy" ]
3
[ "methods: miscellaneous", "minor planets", "asteroids: general", "planet-disc interactions", "circumstellar matter", "planetary systems", "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1980AJ.....85.1122W", "1984Icar...58..109F", "1989Icar...82..402D", "1996Icar..119..261C", "1996Icar..123..180K", "1999A&AS..137..273G", "1999AJ....118.2993G", "1999MNRAS.304..793C", "1999ssd..book.....M", "2003A&A...402..701B", "2004ARA&A..42..549G", "2004ApJ...614..497G", "2005ApJ...635.1348F", "2005Natur.435..466G", "2007ApJ...663..365W", "2007ApJ...666..423Z", "2007MNRAS.382.1823F", "2008ARA&A..46..339W", "2008ApJ...677..630H", "2008Sci...322.1348M", "2009ApJ...693..734C", "2009ApJ...705..314S", "2009Icar..201..381S", "2009MNRAS.399..385B", "2010ApJ...710.1408F", "2010Natur.468.1080M", "2010lyot.confE..42D", "2011ApJ...729..128C", "2011ApJ...732...61Z", "2011ApJ...735..109W", "2011ApJ...739...31L", "2011ApJ...739...36S", "2011MNRAS.418.1043Q", "2012ApJ...761...57B", "2012MNRAS.419.3074M", "2013A&A...558A.121K", "2013ApJ...767..115F", "2013Icar..225...40B", "2014ApJ...787...80H", "2014ApJ...790..146F", "2014IAUS..299..318S", "2014Icar..232..118M", "2014MNRAS.437.3727K", "2014MNRAS.440.3140G", "2014MNRAS.444.3164K", "2014MNRAS.445.4175V", "2014PNAS..11112661M", "2015ApJ...798...83N", "2015ApJ...799...41M", "2015ApJ...800....5M", "2015ApJ...807...44P", "2015MNRAS.446.2059S", "2015MNRAS.448..684S", "2015MNRAS.453.3329P", "2015MNRAS.454.3267F", "2016A&A...587A..56M", "2016ApJ...823..118M", "2016ApJ...825...77D", "2016MNRAS.460L..10B" ]
[ "10.1093/mnrasl/slw143", "10.48550/arXiv.1607.04282" ]
1607
1607.04282_arXiv.txt
Debris disks are circumstellar dust disks, produced by the destructive collisions of planetesimals leftover from the planet formation process \citep{2008ARA&A..46..339W}. There exist a significant number of debris disks with two temperature components \citep{2008ApJ...677..630H}. Modelling suggests that in at least a significant fraction of cases, these two temperature disks harbour two concentric debris rings, with a significant gap between them \citep{2014MNRAS.444.3164K} - somewhat analogous to the asteroid and Kuiper belts of the Solar system. Also by analogy with the Solar system, the gap is often inferred to have been opened by planets scattering away the remnant planetesimals. \citet{2007MNRAS.382.1823F} modelled the gap clearing as caused by multi-planet instabilities \citep{1996Icar..119..261C} producing `Nice model' like clearing of massive planetesimal belts \citep{2005Natur.435..466G}. However, attempts to match such instabilities to observed debris disks suggest they must be rare events overall \citep{2009MNRAS.399..385B}, and thus they are unlikely to be the principle mechanism for gap clearing. This rarity should also apply to the formation of a double ring by a single, eccentric, dynamically unstable planet, as modelled by \citet{2015MNRAS.453.3329P}. The time for a single planet to clear its chaotic zone was considered by \citet{2015ApJ...799...41M} and \citet{2015ApJ...798...83N}. In the case of the observed gaps opened in double debris disk systems, the necessary planet mass is often too large to have escaped detection by direct imaging attempts. This led \citet{2014IAUS..299..318S} to suggest that the observed gaps may be opened by several planets scattering away the remnant planetesimals. Despite some attempts \citep{2007ApJ...666..423Z,2011MNRAS.418.1043Q,2011ApJ...735..109W,2011ApJ...739...31L}, a general theory of the stability of many-planet systems has not yet been developed. Great success, however, has been enjoyed by $N$-body simulations \citep{1996Icar..119..261C,2009Icar..201..381S,2014MNRAS.437.3727K,2015ApJ...807...44P}. Thus, to consider the case where gaps in double debris disks are caused by multiple planets scattering away the planetesimals leftover from the planet formation epoch, we use $N$-body simulations to calculate the clearing time for a given planetary system. By inverting this relation, we recover an equation for the minimal planetary system that must be present in a gap for a system of a given age (figure \ref{fig:schematic}). \begin{figure} \centering {\includegraphics[width=0.49\textwidth,trim = 0 0 0 0, clip]{schematicb.pdf}} \caption{Schematic representation of how the lower limit presented here produces an overall bounded view of the possible planetary system when combined with the upper limits from direct imaging non-detection.} \label{fig:schematic} \end{figure}
Implicit in this approach is the assumption that planets form surrounded by a sea of planetesimals that still retain a significant fraction of their mass. There is some theoretical basis to believe planet formation may be $\sim 50\%$~efficient \citep{2004ApJ...614..497G,2004ARA&A..42..549G}. There is some circumstantial observational evidence of this; the estimated mass of the Oort cloud \citep{2005ApJ...635.1348F,2015MNRAS.454.3267F}~and calculated fraction of small bodies that end up in the Oort cloud \citep[e.g.][]{2013Icar..225...40B,2015MNRAS.446.2059S}~imply the mass of solids scattered by the planets was comparable to the mass of solids in the planets. Similar mass clouds may be commonly present around other stars \citep{2014MNRAS.445.4175V}. Modelling of debris disks also suggests their total mass is comparable to the solid mass of planetary systems \citep{2008ARA&A..46..339W,2011ApJ...739...36S}. The observational evidence does not strongly indicate that the proto-comets were co-spatial with the planets; if future observations fail to find the minimal planetary systems envisioned here, it will be significant evidence that planets do not clear gaps, but rather that planetesimal gaps form because planet formation is $\sim 100\%$~efficient, or that giant planets clear gas gaps that also removes solids \citep[as in][]{2016ApJ...825...77D}. Very recently, \citet{2016ApJ...823..118M} published a study on the maximal planetary system that can fit dynamically between two debris disks. This provides a stronger constraint on older systems, and thus might provide a more stringent upper limit than direct imaging for older systems. This model necessitates a caveat: we have neglected the mass of the planetesimals in our study. If the mass of planetesimals is comparable to, or in excess of, that of the planets, they may cause migration of the planets \citep{1984Icar...58..109F}. \citet{2014Icar..232..118M} published a set of criteria for when planets in a planetesimal disk may start to migrate. If the minimum planetary system predicted by this study is such that migration might occur during the clearing phase, the model presented here may be inappropriate. For the young systems most favourable to direct imaging, and most likely to host double debris disks, the minimum mass will be higher (equation \ref{eq:minmass}), and migration is unlikely to be a concern. For instance, for HD 38206 we inferred at least $1500~m_{\oplus}$ in planets, while a typical A star debris disk is inferred to have a mass of $\sim 10~m_{\oplus}$ \citep{2007ApJ...663..365W}. Consequently, from \citet{2014Icar..232..118M} we expect no migration, which only occurs for $m_p < 3~m_{disk}$. A massive disk would also gravitationally self-excite, spreading the planetesimals \citep{1996Icar..123..180K}, and viscously spreading the small bodies \citep{2013A&A...558A.121K}. This could allow them to encounter secular resonances and be cleared on shorter timescales. As the spreading will depend on the mass and size distribution in the debris, there is no good way to estimate the appropriate timescale.
16
7
1607.04282
The gap between two component debris discs is often taken to be carved by intervening planets scattering away the remnant planetesimals. We employ N-body simulations to determine how the time needed to clear the gap depends on the location of the gap and the mass of the planets. We invert this relation, and provide an equation for the minimum planet mass, and another for the expected number of such planets, that must be present to produce an observed gap for a star of a given age. We show how this can be combined with upper limits on the planetary system from direct imaging non-detections (such as with GPI or SPHERE) to produce approximate knowledge of the planetary system.
false
[ "such planets", "direct imaging non", "planets", "approximate knowledge", "detections", "-", "GPI", "SPHERE", "the minimum planet mass", "the planetary system", "an observed gap", "a given age", "the remnant planetesimals", "upper limits", "the planets", "The gap", "the gap", "the expected number", "two component debris discs", "the mass" ]
9.176764
13.435136
-1
1833875
[ "Glover, Simon C. O.", "Smith, Rowan J." ]
2016MNRAS.462.3011G
[ "CO-dark gas and molecular filaments in Milky Way-type galaxies - II. The temperature distribution of the gas" ]
39
[ "Universität Heidelberg, Zentrum für Astronomie, Institut für Theoretische Astrophysik, Albert-Ueberle-Straße 2, D-69120 Heidelberg, Germany", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, Oxford Road, Manchester M13 9PL, UK" ]
[ "2016ApJ...830...18T", "2017ApJ...834...63R", "2017MNRAS.465.4611P", "2017MNRAS.466..574W", "2017MNRAS.466.3293P", "2017MNRAS.467.4322P", "2017MNRAS.470.2890B", "2017MNRAS.471.2151H", "2018ASPC..517..575D", "2018ApJ...862...49N", "2018ApJ...869...73L", "2018MNRAS.475.1508P", "2018MNRAS.475.2383S", "2018MNRAS.481.1016G", "2018MNRAS.481.1976Z", "2018MNRAS.481.4277F", "2018NewAR..82....1H", "2019ApJ...883..158B", "2019MNRAS.483.4707G", "2019MNRAS.486.4622C", "2020A&A...634A.139W", "2020A&A...641A..17B", "2020A&A...642A..68S", "2020A&A...643A...3S", "2020ApJ...903...30B", "2020MNRAS.492.1465S", "2020MNRAS.499..837K", "2021ApJ...909...56L", "2021MNRAS.507.3952R", "2023A&A...676A..89G", "2023Ap&SS.368...76J", "2023ApJ...944L...8S", "2023FrASS..1072771D", "2023PASA...40...15H", "2023PASA...40...17W", "2024A&A...682A.161S", "2024A&A...685A..30W", "2024MNRAS.tmp.1490D", "2024arXiv240300917K" ]
[ "astronomy" ]
8
[ "astrochemistry", "hydrodynamics", "ISM: clouds", "ISM: molecules", "galaxies: ISM", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1093/mnras/stw1879", "10.48550/arXiv.1607.08253" ]
1607
1607.08253_arXiv.txt
Star formation within our Galaxy appears to occur exclusively within clouds of molecular gas. It is therefore important to understand the physical conditions within these clouds if we want to understand the role they play in star formation. Ideally, we would do this by studying the line emission from molecular hydrogen (H$_{2}$), which is by far the most abundant molecule within these clouds. Unfortunately, however, the temperature within the bulk of these clouds is too low to excite even the lowest accessible rotational transition of H$_{2}$, meaning that H$_{2}$ emission is useful only for tracing atypical regions such as the shock fronts created by protostellar outflows \citep[see e.g.][]{rb01}. Consequently, if we want to learn about the physical conditions within the clouds (temperatures, densities, velocity structure, etc.), we have to rely on information provided by other observational tracers. The most widely used tracer molecule is carbon monoxide (CO). This is the second most abundant molecular species within typical molecular clouds and its lowest rotational transitions are easily excited at the gas temperatures found there. Observations of CO line emission have therefore become one of the main ways in which we study the properties of both Galactic and extragalactic molecular clouds \citep[see e.g.][]{jackson06,hughes13}. However, chemical modelling of clouds has long made clear that CO is not a perfect tracer of H$_{2}$ \citep[see e.g.][]{th85,vdb88}. H$_{2}$ is far more effective than CO at shielding itself against the effects of ultraviolet photodissociation, and as a result the transition from atomic to molecular hydrogen within a cloud illuminated by the interstellar radiation field (ISRF) does not necessarily occur in the same location as the transition from C$^{+}$ to C to CO. Consequently, observations of CO alone miss a significant fraction of the H$_{2}$. This CO-poor or ``CO-dark'' molecular gas has attracted increasing interest in the past few years, prompted by the realization that it may represent a significant fraction of the total molecular gas budget of the Milky Way. Observations \citep{grenier05,allen12,paradis12,langer14} and theoretical studies \citep{wolf10,smith14} suggest that in conditions typical of the local ISM, anywhere between 30\% and 70\% of the total mass of H$_{2}$ may be located in regions that are CO-poor and that are hence not traced by observations of CO emission. Observational constraints on the temperature distribution of the molecular gas in the ISM primarily come from observations of molecules such as CO itself or ammonia (NH$_{3}$) that are only found in large quantities within CO-rich gas. These observations therefore only constrain the temperature of the CO-rich molecular gas. The temperature distribution in the CO-dark phase is much harder to constrain observationally. In cases where there is a bright ultraviolet (UV) background source (e.g.\ an AGN), H$_{2}$ can be observed in absorption via its Lyman-Werner band transitions, allowing one to constrain its temperature, but these observations only probe gas in a very small fraction of the ISM \citep[see e.g.][]{rach09}. Since the CO-dark gas is much less well-shielded against the effects of the ISRF than the CO-rich gas, there are good reasons for believing that its temperature distribution could be significantly different from that in the CO-rich regions, an assumption which seems to be borne out by the high ($T \sim 50$--100~K) H$_{2}$ temperatures recovered from UV absorption line studies. If the temperature of the CO-dark gas does significantly differ from that of the CO-rich gas, then this is important as the temperature of the gas controls its sound-speed, and hence influences its stability against gravitational collapse as well as its response to the effects of supersonic turbulence \citep{pnj97,molina12}. The temperature of the molecular gas also affects the rate at which chemical reactions take place within it and so is an important input into any chemical model of the CO-dark molecular phase. Finally, the temperature of the CO-dark molecular gas also strongly influences the ease with which it can be detected using the atomic fine structure transitions of ionized carbon (C$^{+}$) and atomic oxygen (O). In an effort to better understand the properties of the CO-dark molecular gas in galaxies like our own Milky Way, we have carried out detailed hydrodynamical simulations of a representative portion of the disk of a massive spiral galaxy \citep{smith14}. These simulations combine a state-of-the-art hydrodynamical code with a detailed treatment of the small-scale cooling physics and chemistry, and allow us to study the properties of the molecular phase in unprecedented detail. In this paper, we study in detail the temperature distribution of the molecular gas in these simulations, and assess its detectability using the fine structure lines of C$^{+}$ and O. The structure of our paper is as follows. In Section~\ref{sims}, we briefly summarize the numerical method used for our simulations and the initial conditions that we adopt. (A more lengthy discussion of both can be found in \citealt{smith14}). In Section~\ref{sec:MW}, we investigate the temperature distribution of the H$_{2}$ in our fiducial ``Milky Way'' simulation, and show that the temperatures that we derive for the CO-dark molecular component are in good agreement with previous theoretical predictions. In Section~\ref{sec:change}, we extend the scope of our study to our other three simulations, allowing us to assess the effect of changing the mean gas surface density and the strength of the interstellar radiation field. In Section~\ref{sec:map}, we use a simple approximation to estimate the surface brightness of our simulated galaxies in the \cii 158$\,\mu$m, \oi 63$\,\mu$m and \oi 145$\,\mu$m fine structure lines and investigate how well the emission in these lines traces the CO-dark molecular gas. We discuss our results in Section~\ref{sec:discuss} (including a discussion of some possible caveats) and present our conclusions in Section~\ref{sec:conc}.
\label{sec:discuss} \subsection{The temperature distribution of CO-dark molecular gas} Our simulations demonstrate that the temperature distribution of CO-dark molecular gas is significantly different from the temperature distribution of the molecular gas that is well-traced by CO emission. CO-bright molecular gas has a low mean temperature, which can be as low as 10~K or as high as 30~K, depending on the strength of the ambient ISRF, and also has a narrow temperature distribution. CO-dark molecular gas, on the other hand, has a much higher mean temperature, ranging from 74~K in the Low and Weak simulation to 105~K in the Strong Field simulations. The CO-dark H$_{2}$ also occupies a broader range of temperatures than the CO-rich gas. It is clear from this that temperatures derived using molecular tracers (e.g.\ CO itself or NH$_{3}$) are only useful for describing the behaviour of the CO-bright gas, and not for constraining the temperature of the CO-faint material. In addition, if we compare the temperature distribution of the cold H$\,${\sc i} in our simulations with that of the CO-dark H$_{2}$, we see that although there are more similarities, the two temperature distributions are not the same. H$\,${\sc i} is typically found in slightly warmer gas than CO-dark H$_{2}$, since the conditions that favour H$_{2}$ formation (higher densities, moderate amounts of shielding) also allow the gas to reach lower temperatures than the lower density gas dominated by atomic hydrogen. Therefore, temperature measurements made using H$\,${\sc i} emission or absorption are also only of limited usefulness in constraining the temperature of the CO-dark H$_{2}$. One consequence of this is that the sound speed of the CO-dark molecular gas is significantly lower than that of the cold H$\,${\sc i}. If we compute the mass-weighted mean sound speed of the regions dominated by cold H$\,${\sc i} in our Milky Way simulation (i.e.\ regions with molecular fractions of less than 50\% and temperatures $T < 300$~K), then we find that $c_{\rm s, CNM} = 1.15 \: {\rm km \: s^{-1}}$. This value is only weakly dependent on the value we take for the upper limit on $T$, provided that it is higher than a few hundred K. For comparison, the mass-weighted mean sound speed of the H$_{2}$-dominated regions (i.e.\ those with molecular fractions greater than 50\%) is $c_{\rm s, H_{2}} = 0.47 \: {\rm km \: s^{-1}}$ if we consider all of the H$_{2}$-dominated regions, increasing to $c_{\rm s, dark} = 0.64 \: {\rm km \: s^{-1}}$ if we only include cells in which less than 1\% of the total carbon is locked up in CO. This difference in sound speeds has two important consequences. First, it implies that the Jeans mass in the CO-dark molecular gas will generally be smaller than in the H$\,${\sc i}-dominated regions, meaning that the molecular regions will be more unstable to gravitational collapse. Second, if the turbulent velocity dispersion in the H$\,${\sc i}-dominated gas is similar to that in the CO-dark H$_{2}$, then turbulence will have an easier job creating dense structures in the latter than in the former. For example, we know that for an isothermal, supersonically turbulent gas, the logarithmic density variance scales as \citep{molina12} \begin{equation} \sigma^{2} = \ln \left[1 + b^{2} \left(\frac{\beta}{\beta + 1} \right) {\cal M}^{2} \right], \end{equation} where ${\cal M} \equiv \sigma_{v} / c_{\rm s}$ is the Mach number of the turbulence, $\beta$ is the ratio of the thermal pressure to the magnetic pressure, given by $\beta \equiv 2 c_{\rm s}^{2} / v_{\rm A}^{2}$ where $v_{\rm A}$ is the Alfven velocity, and $b$ is a parameter that depends on the nature of the turbulent forcing \citep{federrath08,federrath10}. In a non-isothermal gas, one recovers a slightly different relationship between $\sigma$ and ${\cal M}$ \citep{nolan15}, but the general point that an increase in ${\cal M}$ leads to an increase in the density variance (i.e.\ an increased number of highly over-dense or under-dense regions) remains valid. Together, these factors mean that it is easier to form gravitationally bound clouds in regions of CO-dark molecular gas than in regions of cold atomic gas. It also implies that simulations of cloud formation in the ISM that do not account for the transition from H$\,${\sc i} to H$_{2}$ will tend to underestimate the formation rate of gravitationally bound clouds. Another important implication of the fact that the CO-dark H$_{2}$ has a relatively broad temperature distribution concerns its detectability using the fine structure lines of \cii and \oino. The excitation rate of \cii due to collisions with H$_{2}$ increases by a factor of ten as we increase the temperature from 30~K to 100~K, and so the warmer the H$_{2}$, the easier it is to trace using \cii emission. Consequently, attempts to map the distribution of CO-dark molecular gas using \cii emission \citep[e.g.][]{langer14} inevitably give us a somewhat biased picture, as they preferentially pick out warmer gas and may not have sufficient sensitivity to detect CO-dark H$_{2}$ in relatively cold regions. In the case of \oino, the excitation rate changes by an even larger factor over the temperature range $30 < T < 100$~K, and so the bias is even greater. \subsection{Caveats} \label{sec:caveats} As previously discussed in \citet{smith14}, there are several caveats that should be taken into account when considering the results of our simulations. First, and most importantly, there is the fact that in the simulations discussed here, we do not account for the effects of either self-gravity or stellar feedback. Although we expect the effects of these processes to cancel out to some extent, with stellar feedback disrupting clouds that would otherwise undergo runaway gravitational collapse, our neglect of both processes nonetheless represents a major simplification in comparison to the real ISM. Note, however, that we would expect both processes to have a much stronger impact on the behaviour of the relatively dense, CO-rich molecular gas than on the comparatively diffuse CO-dark H$_{2}$. A second important simplification in our models is our assumption of a uniform interstellar radiation field. On small scales (e.g.\ close to regions of massive star formation), this assumption clearly breaks down, but on large scales it should be a reasonable approximation \citep{Wolfire03}. Our assumption of a constant metallicity and a constant dust-to-gas ratio also represents a simplification in comparison to the real Milky Way, which has a slight metallicity gradient \citep[see e.g.][]{Luck11}, although for the range of radii we simulate, this corresponds to a change in metallicity of less than a factor of two. Finally, the treatment of CO chemistry used in these simulations (based on \citealt{nl97}) is somewhat approximate and is known to over-produce CO in turbulent clouds in comparison to more sophisticated models \citep{gc12a}. Consequently, we may mis-classify as CO-rich some gas that should actually be CO-dark. However, accounting for this will only broaden the temperature distribution of the CO-dark H$_{2}$, strengthening our main conclusions.
16
7
1607.08253
We investigate the temperature distribution of CO-dark molecular hydrogen (H<SUB>2</SUB>) in a series of disc galaxies simulated using the AREPO moving-mesh code. In conditions similar to those in the Milky Way, we find that H<SUB>2</SUB> has a flat temperature distribution ranging from 10 to 100 K. At T &lt; 30 K, the gas is almost fully molecular and has a high CO content, whereas at T &gt; 30 K, the H<SUB>2</SUB> fraction spans a broader range and the CO content is small, allowing us to classify gas in these two regimes as CO-bright and CO-dark, respectively. The mean sound speed in the CO-dark H<SUB>2</SUB> is c<SUB>s, dark</SUB> = 0.64 km s<SUP>-1</SUP>, significantly lower than the value in the cold atomic gas (c<SUB>s, CNM</SUB> = 1.15 km s<SUP>-1</SUP>), implying that the CO-dark molecular phase is more susceptible to turbulent compression and gravitational collapse than its atomic counterpart. We further show that the temperature of the CO-dark H<SUB>2</SUB> is highly sensitive to the strength of the interstellar radiation field, but that conditions in the CO-bright H<SUB>2</SUB> remain largely unchanged. Finally, we examine the usefulness of the [C II] and [O I] fine-structure lines as tracers of the CO-dark gas. We show that in Milky Way-like conditions, diffuse [C II] emission from this gas should be detectable. However, it is a problematic tracer of this gas, as there is only a weak correlation between the brightness of the emission and the H<SUB>2</SUB> surface density. The situation is even worse for the [O I] line, which shows no correlation with the H<SUB>2</SUB> surface density.
false
[ "CO-dark molecular hydrogen", "CO-dark", "CO", "the CO-dark gas", "gas", "the CO-dark molecular phase", "the CO-dark H<SUB>2</SUB", "the CO-bright H<SUB>2</SUB", "H<SUB>2</SUB", "a high CO content", "gravitational collapse", "the cold atomic gas", "c", "SUB", "turbulent compression", "the CO content", "disc galaxies", "T", "conditions", "the H<SUB>2</SUB> surface density" ]
12.165374
9.414847
-1
12409137
[ "Hu, Huidong", "Liu, Ying D.", "Wang, Rui", "Möstl, Christian", "Yang, Zhongwei" ]
2016ApJ...829...97H
[ "Sun-to-Earth Characteristics of the 2012 July 12 Coronal Mass Ejection and Associated Geo-effectiveness" ]
38
[ "State Key Laboratory of Space Weather, National Space Science Center, Chinese Academy of Sciences, Beijing 100190, China ; University of Chinese Academy of Sciences, No.19A Yuquan Road, Beijing 100049, China;", "State Key Laboratory of Space Weather, National Space Science Center, Chinese Academy of Sciences, Beijing 100190, China ; University of Chinese Academy of Sciences, No.19A Yuquan Road, Beijing 100049, China", "State Key Laboratory of Space Weather, National Space Science Center, Chinese Academy of Sciences, Beijing 100190, China", "Space Research Institute, Austrian Academy of Sciences, A-8042 Graz, Austria ; IGAM-Kanzelhöhe Observatory, Institute of Physics,University of Graz, A-8010 Graz, Austria", "State Key Laboratory of Space Weather, National Space Science Center, Chinese Academy of Sciences, Beijing 100190, China" ]
[ "2017ApJ...834..158L", "2017ApJ...837....4Z", "2017ApJ...840...76H", "2017ApJ...849..112L", "2017SSRv..212.1159M", "2017ScChD..60.1383C", "2017ScChD..60.1466H", "2017SoPh..292...74S", "2017SoPh..292..116H", "2017SoPh..292..142W", "2017SoPh..292..189M", "2018ApJ...857...36Y", "2018ApJ...863...81W", "2018ChJSS..38..665Z", "2018IAUS..335..258G", "2018JGRA..123.7167H", "2018SSRv..214...21R", "2018ZNatA..73..385R", "2019A&A...621A..72A", "2019A&A...626A.122S", "2019ApJ...871....8L", "2019ApJ...878..106H", "2019ApJ...882..122Z", "2019ApJ...883...91F", "2019ApJ...884...90C", "2019ApJS..241...15L", "2019PhLA..38325919K", "2019SpWea..17..357W", "2019sfsw.book..165M", "2019sfsw.book..489R", "2020PhDT........80S", "2021ApJ...919L..30S", "2021PEPS....8...56Z", "2022A&A...667A.133B", "2022AdSpR..70.1641M", "2022ApJ...937...44C", "2023ApJ...955...50P", "2024SpWea..2203715M" ]
[ "astronomy", "physics" ]
7
[ "shock waves", "solar–terrestrial relations", "solar wind", "Sun: coronal mass ejections: CMEs", "Sun: radio radiation", "Astrophysics - Solar and Stellar Astrophysics", "Physics - Space Physics" ]
[ "1971SoPh...17..392F", "1975JGR....80.4204B", "1977SoPh...55..121S", "1984SoPh...90..401B", "1987JGR....9211189L", "1995ApJ...440L.109P", "1998GeoRL..25.2493R", "1998SoPh..183..165L", "1999JGR...104.6899H", "1999JGR...10412515L", "1999JGR...10424739S", "2000GeoRL..27..145G", "2000JGR...105.7707O", "2001ApJ...548L..91G", "2001ApJ...563..381Y", "2001JGR...10629207G", "2001P&SS...49.1445S", "2002A&A...384.1098C", "2002JGRA..107.1142H", "2003ApJ...590..533R", "2003JGRA..108.1361K", "2004JGRA..109.3102H", "2004SoPh..222..329W", "2006ApJ...652..763T", "2006JGRA..111.9208L", "2006JGRA..11111104F", "2007A&A...461.1121C", "2007ApJ...663.1369R", "2007JGRA..112.9103K", "2007SSRv..131..417A", "2008ApJ...675..853S", "2008ApJ...677L.133L", "2008ApJ...689..563L", "2008SSRv..136....5K", "2008SSRv..136...67H", "2008SSRv..136..487B", "2009AnGeo..27.3479L", "2009ApJ...694..707W", "2009GeoRL..36.2102D", "2009JGRA..114.0A22G", "2009JGRA..114.4102M", "2010ApJ...710L..82L", "2010ApJ...715..493L", "2010ApJ...719.1385R", "2010ApJ...722.1762L", "2010GeoRL..3724103M", "2011ApJ...733L..23V", "2012ApJ...746L..15L", "2012ApJ...749...57T", "2012ApJ...750...45H", "2012ApJ...753...21F", "2012ApJ...758...10M", "2012SoPh..275...17L", "2012SoPh..275..229S", "2013ApJ...769...45L", "2013ApJ...777..167D", "2014ApJ...784..144D", "2014ApJ...787..119M", "2014ApJ...789...93C", "2014ApJ...792...49H", "2014JGRA..119.7128S", "2014NatCo...5.3481L", "2015ApJ...807..177G", "2015ApJ...808L..15S", "2015ApJ...809L..34L", "2015ApJ...814...80W", "2015JGRA..120.6101W", "2015NatCo...6.7135M", "2015SoPh..290..891M", "2015SoPh..290.1371M", "2015SoPh..290.2455C", "2015SoPh..290.3343L", "2016ApJS..222...23L", "2016SoPh..291..239G", "2016SoPh..291.1159W" ]
[ "10.3847/0004-637X/829/2/97", "10.48550/arXiv.1607.06287" ]
1607
1607.06287_arXiv.txt
\label{intro} Coronal Mass Ejections (CMEs) are massive expulsions of magnetized plasma from the solar atmosphere. They are called interplanetary CMEs (ICMEs) after traveling into interplanetary space, which are a significant class of triggers of geo-effectiveness. Understanding CME propagation, associated radio bursts, and plasma and magnetic field characteristics in the inner heliosphere is of critical importance for space weather forecasting. Combination of comprehensive remote-sensing and \insitu{} observations is key to these investigations. Previous studies of CME interplanetary propagation categorize CMEs into fast and slow ones by comparing their speed with the average solar wind speed. Slow CMEs experience an acceleration while fast ones decelerate when interacting with the ambient solar wind \citep{sww1999,gll2000}. Combining coronagraph images with \insitu{} measurements, \citet{llr1999,gly2001} obtain empirical models describing propagation of CMEs out to 1 AU. \citet{gly2001} find that a fast CME undergoes a deceleration out to 0.76 AU before moving at a constant speed. These studies are based on coronagraph images with a field of view (FOV) only out to 30 {R$_\sun$} and lack direct measurements between the Sun and Earth. The \textit{Solar Terrestrial Relations Observatory} \citep[\stereo;][]{kkd_stereo2008} with wide-angle imaging sensors is capable of tracking CMEs from the Sun out to the Earth and even beyond. \stereo{} consists of two spacecraft, one leading the Earth (\sta) and the other trailing behind (\stb), which separate by approximately 44 to 45 degrees from each other every year. An identical imaging suite, the Sun-Earth Connection Coronal and Heliospheric Investigation \citep[SECCHI;][]{hmv_secchi2008}, is aboard each spacecraft, which consists of an extreme ultraviolet imager (EUVI), two coronagraphs (COR1 and COR2) and two heliospheric imagers (HI1 and HI2). A radio and plasma wave investigation instrument \citep[SWAVES;][]{bgk_swaves2008} is also mounted, which can detect type II radio bursts associated with CME-driven shocks. Based on coordinated \stereo{} stereoscopic images, a geometric triangulation technique has been developed to track CMEs with no free parameters \citep{ldl2010,ltl2010}. With the triangulation method, it is revealed that fast CMEs impulsively accelerate until even after the X-ray flare maximum, and then rapidly decelerate to a nearly constant speed or gradual deceleration phase \citep{lll2013}, while slow CMEs experience a slow acceleration phase and then travel with a roughly constant speed around the average solar wind level \citep{lhl2016}. A CME could also propagate in a non-radial direction \citep[e.g.,][]{wsw2004SoPh,gmx2009JGRA,mrf2015natco,lpo2015,wld2015apj}, interact with other CMEs \citep[e.g.,][]{lvr2009AnGeo,gyk2001ApJ,llm2012,llk2014NatCo} or co-rotating interaction regions \citep[e.g.,][]{rkf1998,rls2010ApJ,lhl2016}, or rotate in interplanetary space \citep[e.g.,][]{thv2006ApJ,ltl2010,vcn2011ApJ}, which increases difficulty to predict the CME arrival characteristics at the Earth. The \textit{MErcury Surface, Space ENvironment, GEochemistry, and Ranging} \citep[\mes;][]{smg_mes2001} spacecraft provides a great opportunity to study CMEs inside the Earth orbit. \citet{mfk2012} investigate the shocks, flux ropes and interactions between ejecta for a series of CMEs from 2010 July 30 to August 1 with multi-point \insitu{} data from \mes{}, \textit{Venus Express}, \wind{} and \stereo. Radial evolution of a magnetic cloud (MC) is studied based on the data from \mes{} and \stb{} when the two spacecraft were nearly radially aligned in 2011 November \citep{gfr2015}. Using data from the \mes{} magnetometer, \citet{wlp2015} compile 61 ICMEs at Mercury between 2011 and 2014, while \citet{gf2016solphys} identify 119 ICMEs combining the data from \mes{} and \textit{Venus Express}. Interplanetary type II radio bursts emit at the fundamental and/or harmonic of the plasma frequency, which can be applied to determine the radial distance and predict the time of arrival (ToA) of CME-driven shocks at 1 AU using a proper electron density model of the solar wind \citep{rkb2007,llm2008,lll2013,cis2015}. Using the electron density model of \citet{ldb1998}, \citet{lll2013} derive the radial distances of CME-driven shocks from the drift rates of type II bursts, which compare well with those acquired from the geometric triangulation technique based on \stereo{} stereoscopic white-light observations. Besides from the nose of a CME-driven shock, type II radio bursts can also be produced from the flank \citep{ca2002,clm2007}, or a preexisting high-density solar wind structure interacting with the shock \citep[e.g.,][]{rkf1998,rvc2003ApJ,clm2007,fck2012ApJ}. These complicate the estimate of the radial distance of a shock based on an electron density model. Prolonged and enhanced southward magnetic fields associated with ICMEs are important triggers of geomagnetic storms, which depend on the flux-rope orientation and the axial and azimuthal magnetic field components \citep{ywg2001,lhw2015}. A statistical study finds that the tilt angles of ICME flux ropes deduced from \insitu{} force-free flux-rope models are close to the tilt angles of magnetic polarity inversion lines (PILs) in the corresponding solar source regions \citep{may2015}. A Grad-Shafranov (GS) reconstruction technique is capable of estimating the flux-rope orientation, and axial and azimuthal magnetic field components from \insitu{} measurements \citep{hs1999,hs2002}, which has been validated by well separated multi-spacecraft measurements \citep{llh2008,mfm2009}. With the GS reconstruction technique, \citet{lhw2015} find that the flux-rope axial component is a major contributor to the southward magnetic field for the 2015 June 22 geomagnetic storm while the azimuthal component plays a major role in the 2015 March 17 strong geomagnetic storm, which indicates that the flux-rope structure plays an important role in the generation and intensity of geomagnetic activity. On 2012 July 12, a major CME erupted in NOAA AR 11520 associated with an X1.4 class flare that peaked at about 16:49 UT. The active region is the same that caused the 2012 July 23 extreme solar storm \citep{llk2014NatCo}. A strong geomagnetic storm was triggered and reached a minimum D$_\mathrm{st}$ index of $-$127 nT on July 15. The magnetic field configuration and evolution in the solar source region have been investigated with remote-sensing observations and/or three-dimensional magnetohydrodynamics (3D MHD) simulations \citep{cdz2014ApJ,dja2014ApJ,scz2015ApJ,wlw2016}. The CME kinematic evolution has been studied with a 3D MHD model and a semi-empirical drag model by \citet{ssz2014JGRA} and \citet{hz2014ApJ}, respectively. This event is also included as an illustration in statistical analyses of \citet{mah2014ApJ} and \citet{wlp2015}. In this paper we present a comprehensive remote-sensing and \insitu{} analysis of the \thecme. Despite a number of previous studies of this event, our work is unique for reasons given below: (1) a triangulation method based on stereoscopic wide-angle imaging observations is employed to determine the CME kinematics in this study, while the previous studies are limited to single-spacecraft analysis; (2) the stereoscopic imaging results are compared with multi-point \insitu{} measurements at Mercury and the Earth; (3) the frequency drift rate of the type II radio burst associated with the event is analyzed in this work, from which the results are compared with the triangulation analysis; (4) we apply the GS technique to the near-Earth \insitu{} data to reconstruct the flux-rope structure, and the result is compared with the solar source observations to understand the association of the flux-rope properties with the generation of the geomagnetic storm. As far as we know, the comparison between stereoscopic wide-angle imaging observations and multi-point \insitu{} measurements, which can give crucial information on CME kinematics and structure, is still lacking. The comparison of wide-angle imaging observations with long-duration interplanetary type II radio bursts can also yield important knowledge of CME kinematics as well as source regions of the type II bursts, which has not been sufficiently studied. The examination of the CME magnetic structure and its comparison with the analysis of the solar source region are key to understanding how the CME structure/solar source signature is connected with geomagnetic activity. We describe the solar source signatures in Section \ref{source} and propagation characteristics in interplanetary space in Section \ref{prop}. Section \ref{earth} examines the properties of the flux rope and the associated geo-effectiveness. These results are summarized and discussed in Section \ref{discuss}. This work illustrates an end-to-end study of the space weather chain from the Sun to Earth, highlighting the importance of multi-spacecraft remote-sensing and \insitu{} observations in understanding the Sun-to-Earth characteristics and geo-effectiveness of CMEs.
\label{discuss} We have performed a comprehensive analysis of the \thecme, covering the solar source observations by \sdo, the stereoscopic remote-sensing observations from \stereo, the magnetic field signatures at \mes, and the type II radio burst and \insitu{} characteristics observed by \wind. A GS reconstruction is employed to understand the ICME structure and how the structure controls the geomagnetic activity. These results are summarized and discussed below, which illustrate the importance of multi-spacecraft remote-sensing and \insitu{} observations for understanding the physical processes of CME propagation and space weather forecasting. \begin{enumerate} \item This study compares \stereo{} stereoscopic wide-angle imaging observations with multi-point \insitu{} measurements at Mercury and the Earth, which places a strong constraint on CME Sun-to-Earth propagation. The CME kinematics determined from the triangulation technique predicts well the shock arrival time at \mes{} with an error of about 1 hour in this study. A reasonable accuracy is also obtained when we compare the predicted arrival time and speed with the \insitu{} measurements near the Earth. From the Sun to Earth, the CME undergoes an impulsive acceleration, a rapid deceleration, and then a gradual deceleration out to 1 AU, which agrees with the conclusions of \citet{lll2013,lhl2016}. We find that the rapid deceleration ceases at $\sim$50 R$_\sun$ before the CME reached \mes{} ($\sim$0.47 AU), which is different from the coronagraph findings in \citet{gly2001}. Combining this case with the three events in \citet{lll2013}, we suggest that fast CMEs are likely to terminate deceleration before reaching 100 R$_\sun$ from the Sun, which should be noticed in kinematics models of fast CMEs. This work once again proves the reliability of the stereoscopic triangulation technique in determining the kinematics of earthward CMEs in the inner heliosphere. \item Our comparison between wide-angle imaging observations and the interplanetary type II radio burst indicates important clues on the source region of the type II burst. The study reveals that the consistency between the shock distance derived from the type II radio burst and the CME kinematics determined by the triangulation method requires an unusually high solar wind electron density. The type II radio burst was probably produced from a high-density region. This discrepancy and the slow increasing distance of the CME-driven shock derived from the type II radio burst can be explained by the shock interacting with nearby streamers \citep[e.g.,][]{rvc2003ApJ,clm2007,fck2012ApJ}. Another possibility is that the type II radio burst was generated by the shock flank region with lower heights \citep{ca2002,kcr2003JGRA}. Because only one spacecraft detected the type II radio burst, it is not possible to determine the position of the source region by the radio triangulation method of \citet{mrb2015SoPh}. This result implies uncertainties in the determination of CME kinematics using type II radio bursts alone. \item We reconstruct the ICME structure near the Earth in order to connect it with the solar source observations and to understand how the structure contributes to the geomagnetic storm. The flux-rope inclination angle and chirality at 1 AU are consistent with those inferred from the observations of the solar source. It is worth noting that, however, a flux rope may rotate in the corona and interplanetary space \citep{ltl2010,vcn2011ApJ}. The prolonged southward magnetic field near the Earth is mainly from the axial component of the largely southward inclined flux rope. The axial magnetic field component of the flux rope is about two times as large as the azimuthal component as revealed by the GS reconstruction. If the flux rope had not been inclined southward to the large angle in this event, the strength and duration of the southward magnetic field would be much smaller. \citet{lhw2015} reported an event with the azimuthal magnetic field component much larger than the axial component, which suggests that a southward orientation is not a necessity for a strong southward magnetic field. These results indicate the importance of predicting both the flux-rope orientation and magnetic field components in geomagnetic activity forecasting. \end{enumerate}
16
7
1607.06287
We analyze multi-spacecraft observations associated with the 2012 July 12 coronal mass ejection (CME), covering the source region on the Sun from the Solar Dynamics Observatory, stereoscopic imaging observations from the Solar Terrestrial Relations Observatory (STEREO), magnetic field characteristics from Mercury Surface, Space Environment, Geochemistry, and Ranging (MESSENGER), and type II radio burst and in situ measurements from Wind. A triangulation method based on STEREO stereoscopic observations is employed to determine the kinematics of the CME, and the outcome is compared with the results derived from the type II radio burst using a solar wind electron density model. A Grad-Shafranov technique is applied to Wind in situ data to reconstruct the flux-rope structure and compare it with the observations of the solar source region, which helps in understanding the geo-effectiveness associated with the CME structure. Our conclusions are as follows: (1) the CME undergoes an impulsive acceleration, a rapid deceleration before reaching MESSENGER, and then a gradual deceleration out to 1 au, which should be considered in CME kinematics models; (2) the type II radio burst was probably produced from a high-density interaction region between the CME-driven shock and a nearby streamer or from the shock flank with lower heights, which implies uncertainties in the determination of CME kinematics using solely type II radio bursts; (3) the flux-rope orientation and chirality deduced from in situ reconstructions at Wind agree with those obtained from solar source observations; (4) the prolonged southward magnetic field near the Earth is mainly from the axial component of the largely southward inclined flux rope, which indicates the importance of predicting both the flux-rope orientation and magnetic field components in geomagnetic activity forecasting.
false
[ "CME kinematics models", "CME kinematics", "solar source observations", "CME", "magnetic field components", "STEREO stereoscopic observations", "stereoscopic imaging observations", "magnetic field characteristics", "Wind", "the type II radio burst", "a solar wind electron density model", "multi-spacecraft observations", "geomagnetic activity forecasting", "the type II radio", "situ measurements", "situ data", "the solar source region", "lower heights", "II radio burst", "the CME structure" ]
13.309276
16.240068
2
12464919
[ "Kallosh, Renata", "Linde, Andrei", "Roest, Diederik", "Wrase, Timm" ]
2016JCAP...11..046K
[ "Sneutrino Inflation with α-attractors" ]
29
[ "Stanford Institute of Theoretical Physics and Department of Physics, Stanford University, Stanford, 94305 CA, U.S.A.", "Stanford Institute of Theoretical Physics and Department of Physics, Stanford University, Stanford, 94305 CA, U.S.A.", "Van Swinderen Institute for Particle Physics and Gravity, University of Groningen, Nijenborgh 4, 9747 AG Groningen, The Netherlands ; Theoretical Physics Department, CERN, CH-1211 Geneva 23, Switzerland;", "Institute for Theoretical Physics, TU Wien, A-1040 Vienna, Austria" ]
[ "2016JCAP...11..002L", "2016JCAP...11..028M", "2016PhRvD..94l3503R", "2016arXiv161208878R", "2017EPJC...77..469C", "2017JCAP...02..028L", "2017JCAP...06..027D", "2017JHEP...04..144K", "2017JHEP...06..109G", "2017NuPhB.916..688B", "2017PhLB..773..179N", "2017PhRvD..95l3518N", "2018ApJ...863..133N", "2018IJMPD..2750076R", "2018JCAP...06..011H", "2018PhRvD..97f3525D", "2019AnPhy.41167991M", "2019JCAP...08..002L", "2019JCAP...09..040E", "2019JHEP...04..172B", "2019JHEP...05..211K", "2020Ap&SS.365...97K", "2020ApJ...890...58R", "2020IJMPD..2950077S", "2021JCAP...09..017G", "2021JCAP...11..022M", "2021PhRvD.104h3015G", "2022JCAP...07..042M", "2024arXiv240313882H" ]
[ "astronomy", "physics" ]
5
[ "High Energy Physics - Theory", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology" ]
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[ "10.1088/1475-7516/2016/11/046", "10.48550/arXiv.1607.08854" ]
1607
1607.08854_arXiv.txt
\setcounter{equation}{0} Inflation in the early universe is well described by the general class of $\alpha$-attractor models. These can be constructed in supergravity as F-term models \cite{Kallosh:2013yoa} or as D-term models \cite{Ferrara:2013rsa}, where the first $\alpha$-dependent generalization of the Starobinsky potential was introduced. A robust feature of these models is the spectral index $n_s\approx 1-{2\over N}$, which is in perfect agreement with the observed value of $n_s$ for the number of $e$-foldings $50 \lesssim N \lesssim 60$, as one can see in Fig. 12 in \cite{Ade:2015lrj}. The level of gravity waves in $\alpha$-attractors is given by $r\approx {12\alpha / N^2}$, so that $r\approx 4 \alpha \cdot 10^{-3}$ at $N\approx 55$. It is directly related to the curvature of the moduli space ${\cal R}_K= -{2/ (3 \alpha)}$ which in ${\cal N}=1$ supergravity is a free parameter \cite{Ferrara:2013rsa}. From the current Planck-BICEP-Keck constraint $r< 0.07$ \cite{Ade:2015lrj,Array:2015xqh} it follows that $\alpha \lesssim 17$. Future B-mode experiments will constrain this parameter further. More precise constraints on $\alpha$ depend on the details of the models and on the mechanism of reheating, see e.g. \cite{Ellis:2015pla,Kallosh:2016gqp}. Inflation in supergravity is often described using two chiral superfields, the inflaton $T$ and the stabilizer $S$, as first suggested in \cite{Kawasaki:2000yn} for quadratic chaotic inflation, and developed for a general class of models in \cite{Kallosh:2010ug}. There are 4 real scalar fields in these models, one of which is the relatively light inflaton and 3 others vanish during inflation. There are situations where some of these 3 fields may acquire tachyonic masses, which may lead to instability. Fortunately, there are well established methods either to control these instabilities by adding stabilizing terms to the \K\ potential \cite{Kallosh:2010ug}, or to make these 3 extra fields disappear in the context of the theory of nilpotent orthogonal superfields \cite{Ferrara:2015tyn,Carrasco:2015iij,Dall'Agata:2015lek}. In realistic theories of elementary particles one may encounter more than two chiral superfields, which again requires an investigation of stability during and after inflation. There are some general conditions which allow to avoid tachyonic instability for matter fields \cite{Kallosh:2016ndd}. These conditions are easy to satisfy if the fields $S$ and $T$ belong to a hidden sector and do not have direct interactions with matter fields in the \K\ potential and in the superpotential \cite{Kallosh:2016ndd}. However, some popular inflationary models do not belong to this class. For example, Higgs inflation in supergravity is supposed to be driven by the Higgs field, which does not belong to the hidden sector. This may result in instabilities during and after inflation. However, all of these instabilities can be successfully controlled \cite{Ferrara:2010yw}. In this paper we will study sneutrino inflation \cite{Murayama:1992ua,Nakayama:2016gvg}, which is yet another scenario where the stability conditions found in \cite{Kallosh:2016ndd} may be violated \cite{Nakayama:2016gvg}, which typically results in the post-inflationary tachyonic instability. We should note from the very beginning that it is not the goal of our paper to formulate a fully consistent cosmological model of sneutrino inflation complemented by the theory of reheating, leptogenesis, SUSY breaking and many other aspects of this scenario closely related to particle phenomenology. These aspects have been studied in multiple papers on sneutrino inflation, see e.g. a discussion of reheating in \cite{Nakayama:2016gvg}, leptogenesis in \cite{Bjorkeroth:2016qsk}, and SUSY breaking in \cite{Linde:2016bcz}. Here we will limit ourselves to the investigation of inflationary aspects of this scenario as compared to the latest Planck data \cite{Ade:2015lrj,Array:2015xqh}. In particular, we will show that one can improve consistency of sneutrino inflation with the latest Planck data by embedding it into the theory of $\alpha$-attractors. We will also develop a new approach to matter field stabilization in this scenario, which renders the post-inflationary tachyonic instability in this scenario harmless. The latest version of the sneutrino inflation \cite{Nakayama:2016gvg} refers to an extension of the model \cite{Kallosh:2010ug} with the neutrino sector, such that the fermions in the inflaton and stabilizer superfields provide the right-handed partners to the Standard Model neutrinos. Then one generates light neutrino masses via the seesaw mechanism \cite{seesaw}. An interesting observation of \cite{Murayama:1992ua,Nakayama:2016gvg} is that in large field models of inflation the inflaton mass is of order $M\sim 10^{13}$ GeV, which is close to the right-handed neutrino mass scale suggested by the seesaw mechanism and the neutrino oscillation experiments. Explicitly, the left handed neutrino masses come from integrating out the heavy fermions with mass $M$ from the inflationary sector, so that \be m_{\rm light}^\nu \sim {v^2\over M} \sim {10^4\over 10^{13} }{\rm \ GeV} \sim {\rm \ eV}\,, \ee where $v$ is the Higgs vev. The left-handed neutrinos are many orders of magnitude lighter than all the other particles in the Standard Model and the seesaw mechanism is the only known way to explain this, but it does not give a reason for the existence of the heavy fermions of mass $\sim 10^{13}$ GeV. Inflation seems to naturally supply these heavy fields at the right scale. In particular, the general supergravity constructions \cite{Kawasaki:2000yn,Kallosh:2010ug} necessarily involve two superfields (the inflaton and the stabilizer fields) and thus provide two right-handed neutrinos, allowing for two massive and one massless neutrino. The model of sneutrino inflation proposed in \cite{Nakayama:2016gvg} is based on the supergravity implementation of the generalized natural inflation scenario proposed in \cite{Kallosh:2014vja} with the superpotential \footnote{The authors of \cite{Nakayama:2016gvg} call it multi-natural inflation \cite{Czerny:2014wza}, but the supergravity versions of the multi-natural inflation models proposed in \cite{Czerny:2014wza} are quite different; they involve axion superfields without the stabilizer superfield $S$.} \be W= M S \sum_n {g_n \over n} f \sin {\lp n \varphi\over f\rp} \,. \ee The simple natural inflationary model with a single frequency is disfavored by the recent Planck data \cite{Ade:2015lrj,Array:2015xqh}. Moreover, when adding a second frequency term in the superpotential to fit the data, it was observed in \cite{Nakayama:2016gvg} that its coefficient must have a complex coefficient $g_2= Ce^{i\theta}$. Models with complex parameters in the superpotential are somewhat complicated, as shown in \cite{Kallosh:2014xwa}. To fit the future data on $n_s$ and especially $r$ will require a fine tuned choice of $f, C, \theta$ parameters. This is certainly possible but will become more difficult when the bound on $r$ will decrease further, as one can see from Fig. 2 in \cite{Nakayama:2016gvg} where various numerical examples are plotted. The important feature of the multi-frequency natural inflation models which allows them to be used for sneutrino inflation is the unbroken discrete shift symmetry $\vp \rightarrow \vp + 2\pi f$. This symmetry allows one to keep the inflationary predictions intact when 4 additional superfields, 3 leptons $L^a$, $a=1,2,3$ and the Higgs $H^u$ are added to the model via the following superpotential \cite{Nakayama:2016gvg} \be W= p_a L^a H^u \sum_n {g'_n \over n} f \sin {\lp n \varphi\over f\rp}\,. \ee This is invariant under $\vp \rightarrow \vp + 2\pi f$, thus even at large $\vp$ the masses of the lepton and Higgs supermultiplets are not blowing up since $\sin {\lp n \varphi / f\rp}$ is restricted for any value of $\vp$. However, the $\alpha$-attractor models do not have an unbroken discrete shift symmetry; they have instead a symmetry associated with the hyperbolic geometry. Will these models provide a framework for sneutrino inflation models? The answer to this question is positive, as we will argue below. In hyperbolic $\alpha$-attractor models, the leptons and the Higgs field interact with the disk variable $Z=\tanh {\lp \Phi/ \sqrt 6 \alpha\rp} $. In the simplest case we have \be W = p_a \, L^a \, H^u \, Z = p_a \, L^a \, H^u \, \tanh {\lp \Phi \over \sqrt {6 \alpha}\rp} \,, \ee analogous to a single Fourier term in natural inflation. Inflation occurs when $Z$ is very close to the boundary of the moduli space at $Z \bar Z \approx 1$, which corresponds to $\vp \gg \sqrt 6 \alpha$ in terms of the canonically normalized inflaton field $\vp \equiv$Re$(\Phi)$. For a \K\ potential that describes hyperbolic geometry supplemented with a generic superpotential that is non-singular at the boundary of the moduli space, the inflaton potential $V(\vp)$ becomes flat. In this limit, the masses of all particles at large values of the inflaton field $\vp$ approach exponentially fast their constant values corresponding to the limit $Z\to 1$ or $\vp \to \infty$, and all coupling constants of the field $\vp$ to the matter fields become exponentially small in the limit $\vp \gg \sqrt 6 \alpha$ \cite{Kallosh:2016ndd,Kallosh:2016gqp}. Thus the hyperbolic geometry of $\alpha$-attractors protects the masses of leptons and Higgs fields during inflation, and suggests inflationary models which are more flexible with regard to future data and have an elegant geometric interpretation. Instead of periodic functions of the inflaton $\sin {\lp n \varphi / f\rp}$ interacting with matter multiplets, we have a natural geometric variable of the moduli space with a boundary, the coordinate of the Poincar\'e disk $Z= \tanh {\lp \Phi / \sqrt{6 \alpha}\rp}$ where $Z\bar Z<1$. An interpretation of this moduli space in terms of Escher's picture Circle Limit IV was proposed in \cite{Kallosh:2015zsa}, where it was shown that the radius square of this disk is given by $3\alpha$. Thus the trigonometric restriction on couplings to matter is replaced by the hyperbolic one: \be \sin {\lp n \varphi\over f\rp} \leq 1 \, \qquad \text{is replaced by} \qquad Z\bar Z < 1\,, \ee while the role of the axion decay constant $f$ is played by the curvature parameter $\alpha$. This suggests that one may consistently implement sneutrino inflation in the context of the theory of $\alpha$ attractors. An advantage of implementing inflation in the theory of $\alpha$-attractors becomes obvious when one recalls that the consistency of the model of Ref. \cite{Nakayama:2016gvg} with the Planck data requires the fine-tuning of 2 parameters, $f$, and $\theta$, see Fig. 2 in \cite{Nakayama:2016gvg}. Meanwhile the predictions of $\alpha$-attractors are compatible with observational data for all $\alpha \lesssim O(10)$, without any fine-tuning. To verify the consistency of the sneutrino inflation scenario one should check the stability of the system with respect to the generation of large vevs of the many scalar fields involved. Indeed, in this scenario the interactions of the inflaton field with matter do not satisfy the stability conditions formulated in \cite{Kallosh:2016ndd}. It was argued in \cite{Nakayama:2016gvg} that the stability of the inflationary regime in this scenario can be achieved by adding stabilizing terms to the \K\ potential. We confirm this conclusion, but find that unless the coefficients in front of these terms are incredibly large, the tachyonic instability does appear after the end of inflation. Adding the stabilizing terms to the \K\ potential does not help here; moreover, once the instability starts, these terms make the development of the instability unpredictable. Fortunately, we found a simpler way to achieve matter field stabilization during inflation, without introducing higher order terms in the \K\ potential. Even with this novel stabilization mechanism, the tachyonic instability does appear after inflation, but we found that in our scenario this instability is transient and harmless. It is just a part of a post-inflationary tachyonic preheating \cite{Felder:2000hj}, similar to the waterfall regime at the end of the hybrid inflation scenario \cite{Linde:1991km}. Our results may have implications for general models of multi-field inflation, going beyond sneutrino inflation. In multi-field models, the evolution of scalar fields may go along a very complicated path in a random potential, which may render long and stable large-field inflationary trajectories unlikely, see e.g. \cite{Aazami:2005jf,Marsh:2013qca,Freivogel:2016kxc}. The situation improves if one of these fields belongs to the class of $\alpha$ attractors, which requires it to have a singular kinetic term at some point in the moduli space \cite{Kallosh:2013daa}. Indeed, it is easier to avoid instabilities in a vicinity of a single point. Then, upon transition to canonical variables, the potential acquires an infinitely long stabilized flat inflaton direction. At large values of the inflaton field along this direction, it essentially decouples from all matter fields, making them irrelevant for the development of inflation \cite{Kallosh:2016gqp}. When inflation along this direction ends, all other fields may fall down in a non-inflationary manner without affecting the main inflationary predictions and creating the problems described in \cite{Aazami:2005jf,Marsh:2013qca,Freivogel:2016kxc}. The main remaining problem is to implement these ideas in supergravity and make sure that these fields do not fall to a collapsing vacuum state with a negative vacuum energy. This problem does not appear, if the theory has a stable supersymmetric Minkowski vacuum in which it settles at the end of inflation. We will show that this is indeed the case in the sneutrino inflation model that we analyze. Of course, this is not the end of the story, since our vacuum is not supersymmetric, but this is not a real problem either, because one can break supersymmetry and uplift the vacuum in a way that does not affect the inflationary evolution, see e.g. \cite{Linde:2016bcz} and references therein. This suggests that the theory of cosmological attractors in combination with vacuum stabilization may significantly simplify the construction of realistic multifield inflationary models in supergravity. The outline of this paper is as follows. In section 2 we will provide a review of $\alpha$-attractor models. These will be coupled to the neutrino sector in section 3. The stability during and after inflation and the role of tachyonic directions is discussed in section 4. We close with a discussion and present our conclusions in section 5.
In this paper we studied how to use the inflationary $\alpha$-attractor models to develop the idea of the minimal sneutrino inflation of the kind proposed in \cite{Nakayama:2016gvg}, which was based on supergravity implementation of the natural inflation scenario \cite{Kallosh:2014vja} with a discrete shift symmetry. In \cite{Nakayama:2016gvg} the coupling of the inflaton to leptons and Higgs was arranged to depend on $\sin (n\vp/f)\leq 1$, so that it would never become very strong, even when the canonically normalized inflaton field $\vp$ takes super-Planckian values. We argued that the same goal is easily reached in the context of the $\alpha$-attractor models since the coupling of the inflaton to leptons and Higgs is arranged using the natural geometric Poincar\'e disk variable $Z=\tanh {\lp\vp / \sqrt{6\alpha}\rp}<1$. This stabilizes the behavior of masses and coupling constants during inflation, making the inflaton coupling to matter exponentially small at large values of the inflaton field \cite{Kallosh:2016gqp}. Yet another advantage of $\alpha$-attractors is that they provide one of the best fits to the existing observational data \cite{Ade:2015lrj}. The attractive feature of the minimal sneutrino inflation proposed in \cite{Nakayama:2016gvg}, that is also a feature of our $\alpha$-attractor based models, is the possibility to realize the seesaw mechanism. Two heavy right-handed neutrinos are residing in the inflationary sector of the theory (inflaton and stabilizer superfields), and have characteristic couplings associated with the scale of inflation $M\sim 10^{-5}M_{Pl} \sim 10^{ 13}$ GeV. The light neutrinos then get the masses \be m_{ab} ^\nu = {v^2 \sin^2 \beta \over M} \lp p_a p_b +q_a q_b\rp \,. \ee For $p_a,q_a\sim 1$ and $v^2\sim 10^{4} $ GeV we find the scale of the light neutrinos to be of the order $10^{4} \cdot 10^{-13}$ GeV which is of the order eV, where they should be. In order to make this model minimal and to be able to describe the physics of neutrinos properly with regards to the MNS matrix, one has to make the inflaton sector interact directly with leptons and the Higgs. This takes the inflationary sector out of the hidden sector, which may lead to stability issues. It was recently shown in \cite{Kallosh:2016ndd} that when inflation is in the hidden sector, adding matter causes no tachyonic instabilities under certain conditions which are not very restrictive. It is particularly important for this purpose to avoid a direct interaction between matter fields and the stabilizer $S$, since this interaction may lead to tachyons. Indeed, in agreement with \cite{Kallosh:2016ndd}, the negative terms in the mass eigenvalues in eq. \eqref{eigen2} are due to the coupling of the leptons and Higgs field to the stabilizer $S$. There are several different mechanisms which can be used to stabilize all fields during inflation, so that the presence of matter does not affect the inflationary evolution. The method proposed in \cite{Nakayama:2016gvg} is based on adding stabilizing higher order terms $- \zeta S\bar S L^{2} \bar L^{2}$ to the \K\ potential. We found that such terms can eliminate the tachyonic instability during inflation, but they become inefficient at small values of the inflaton field during the post-inflationary evolution. Once the tachyonic instability develops, it may bring the fields dangerously close to the boundary of the moduli space, which makes investigation unreliable. We developed an alternative, simpler method of stabilization. Our mechanism does not eliminate the tachyonic instability, but it keeps it under control and renders it transient and harmless. This tachyonic instability becomes a part of the post-inflationary tachyonic preheating \cite{Felder:2000hj}. A detailed description of this nonperturbative process in sneutrino inflation requires lattice simulations similar to those performed in \cite{Felder:2000hj}, but we do not expect any surprises here: tachyonic preheating tends to stimulate the energy transfer from the inflaton to the matter fields, thus making reheating faster and even more efficient than the process considered in \cite{Nakayama:2016gvg}. However, the process studied in \cite{Nakayama:2016gvg} is already very efficient because the interactions of the inflaton field with other matter fields is rather strong, which leads to a very high reheating temperature $T \sim 10^{14} - 10^{15}$ GeV. The tachyonic preheating makes the whole process much more involved than the process studied in \cite{Nakayama:2016gvg} which ignores the tachyonic instability, but this should not substantially alter the resulting estimate of the reheating temperature. In supergravity, a high reheating temperature is not necessarily a good thing, because a reheating temperature above $10^{8} $ GeV may lead to the cosmological gravitino problem \cite{Ellis:1982yb,GRAV,KKM}. One may solve this problem by considering a superheavy gravitino with mass $m_{3/2} \gtrsim 10^{2}$ TeV \cite{Nakayama:2016gvg,Dudas:2012wi,Ellis:2015jpg}. However, in that case one must get rid of the excess of LSP produced by gravitino decay, e.g. by introducing R-parity violation \cite{Nakayama:2016gvg}. Thus there is a price to pay for the high reheating temperature, which appears in sneutrino inflation, in Higgs inflation, as well as in any other inflationary model in supergravity in which the inflaton does not belong to the hidden sector. On a positive side, a high reheating temperature requires a greater number $N$ of e-foldings of inflation, which slightly increases the spectral index $n_{s} = 1-2/N$ for $\alpha$-attractors. This may be an advantage of this class of models since it may further improve the compatibility of these models, as well as of the Starobinsky model, with the latest observational data \cite{Ade:2015lrj,Ellis:2015pla,Kallosh:2016gqp}. Independently of all of these issues, we believe that the investigation of the tachyonic instability in this model provides an interesting example of a very unusual dynamical behavior where an instability may lead to the generation of scalar fields of nearly Planckian amplitude, which, however, tend to disappear after the subsequent cosmological evolution. Thus, instead of trying to avoid the tachyonic instability, one may try to learn whether one can use the new possibilities constructively. A similar waterfall tachyonic instability was first discovered in the context of the hybrid inflation \cite{Linde:1991km}. This effect resulted in the creation of superheavy cosmic strings, which ruled out some of the most popular versions of hybrid inflation in supergravity. In the new set of models discussed above, this waterfall regime can end up in a stable supersymmetric Minkowski vacuum without producing any undesirable cosmological defects. One can break supersymmetry in this vacuum and uplift it to dS with a small cosmological constant without affecting investigation of inflationary evolution \cite{Linde:2016bcz}. As we already argued in the Introduction, these results imply that a combination of cosmological attractors and vacuum stabilization in supergravity may significantly simplify the construction of realistic inflationary models involving a large number of scalar fields.
16
7
1607.08854
Sneutrino inflation employs the fermionic partners of the inflaton and stabilizer field as right-handed neutrinos to realize the seesaw mechanism for light neutrino masses. We show that one can improve the latest version of this scenario and its consistency with the Planck data by embedding it in the theory of cosmological α-attractors.
false
[ "light neutrino masses", "attractors", "cosmological α", "-", "the seesaw mechanism", "Sneutrino inflation", "right-handed neutrinos", "the Planck data", "the inflaton and stabilizer field", "the fermionic partners", "the latest version", "the theory", "its consistency", "this scenario", "We", "it", "one" ]
9.940096
-1.515975
86
12406287
[ "Ioannidis, P.", "Schmitt, J. H. M. M." ]
2016A&A...594A..41I
[ "Glimpses of stellar surfaces. I. Spot evolution and differential rotation of the planet host star Kepler-210" ]
12
[ "Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112, 21029, Hamburg, Germany", "Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112, 21029, Hamburg, Germany" ]
[ "2018ApJ...865..142B", "2018MNRAS.474.5534O", "2018exha.book.....P", "2020A&A...641A.158M", "2020A&A...644A..26I", "2020ARep...64..556A", "2020ApJ...901...14B", "2021A&A...649A..17D", "2021MNRAS.501.1878X", "2021MNRAS.508.5687H", "2022A&A...657A..37M", "2022ApJ...940..132V" ]
[ "astronomy" ]
2
[ "stars: activity", "starspots", "stars: rotation", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1051/0004-6361/201628491", "10.48550/arXiv.1607.08065" ]
1607
1607.08065_arXiv.txt
\begin{figure*}[t] \includegraphics[width=\linewidth]{kepler210_lc} \caption{The complete light curve of the Kepler-210 ($\sim$1$\,$400 days), normalized and with the planetary transits excluded (see text for details). } \label{fig:lght} \end{figure*} The nature of the magnetic dynamo and the role of differential rotation in dwarf main-sequence stars is not yet fully understood. While these types of stars often exhibit substantial photometric variability, the amplitude of their surface differential rotation is expected to be small \citepads{2008JPhCS.118a2029K}. The inability of the small or absent differential rotation to organize the global magnetic field of those stars may result in the absence of activity cycles as observed in the Sun \citepads{2006A&A...446.1027C}. As a result, the study of photospherically active stars and the measurement of their differential rotation rates is an important piece of information for all dynamo theories. The rotational periods of stars can be measured with a variety of techniques including monitoring of the {intensity variations of the cores of the Ca~H+K lines}, spectral line broadening (for cases with known stellar radius and inclination), and the analysis of pseudo-periodic photometric modulations as a result of surface inhomogeneities in the form of photospheric activity (spots). The successful operation of the space missions {\it CoRoT} \citepads{2006ESASP1306...33B} and {\it Kepler} \citepads{2010Sci...327..977B}, in combination with their {high photometric} accuracy and long, non-interrupted observations has revolutionized the studies of stellar rotation with period measurements being available for a very large number of field stars (\citeads{2013A&A...560A...4R}, \citeads{2014ApJS..211...24M}). Furthermore, the analysis of the photometric light curves of {\it CoRoT} and {\it Kepler,} using a variety of techniques, including power spectrum analysis \citepads{2013A&A...557A..11R} and spot modeling (\citeads{2009A&A...506..263F}, \citeads{2009A&A...508..901H}, \citeads{2011A&A...532A..81F}, \citeads{2012A&A...547A..37B}), makes it possible to measure the stellar differential rotation and other physical characteristics of star spots. In this paper we use the light curve phenomenology of Kepler-210 (i.e., the variations of the light curve between the stellar rotations) in combination with spot modeling and power spectrum analysis to study its photospheric activity. In the first part of Sect.~\ref{sect:data} we describe the data and the properties of Kepler-210. In the second part of this section we present the details of the spot model used in our analysis and explain our choices regarding the number of free parameters in our model. In Sect.~\ref{sect:sp_pos}, we show the results of our combined analysis and, in Sect.~\ref{sect:ph_int}, we attempt to {provide a physical interpretation of our results}. Finally, we conclude with a summary in Sect.~\ref{sect:conc}.
\label{sect:conc} Using the phenomenology of the Kepler-210 light curve in combination with the results of a five-spot model, we study the behavior of the spotted areas on the star (i.e., their relative periods and the longitudes at which they appear) and their changes in time (see Fig.~\ref{fig:res}). Based on the spot phenomenology we identify six different ``spot seasons'' and demonstrate that there are differences in the dominant periods of the L-S periodograms corresponding to each season (see Fig.~\ref{fig:ls-seasons}). Additionally we show that the seasons also manifest themselves as differences in the correlation between the corresponding parts of the light curve (see Fig.~\ref{fig:croscor}). According to Fig.~\ref{fig:ls-seasons}, the relative period of spots in the subsequent seasons appears to change in the same fashion as the relative rotational period of sunspots during the solar cycle, i.e., the relative starspot period appears to change from lower to higher values between seasons S2 and S3, while it diminishes gradually from season S4 until the end of the {\it Kepler} observations. A common characteristic between all seasons, with the exception of seasons S3 and S4, is the persistent appearance of the spots in a specific longitude range of the star. Assuming solar-like differential rotation we show that the value of the strength of the differential rotation $\alpha$ {ought to be similar or lower to that of the Sun under the hypothesis that the spots with the higher periods have latitudes in the range of between $25\degree \lesssim\phi_\mathrm{max}\lesssim40\degree$ (see Fig.~\ref{fig:difrot}).} Furthermore, we estimate the spot life times using the radii of the spots that were computed with our model fit. As a result, we conclude that the spot life times of Kepler-210 vary between $\sim$60~days and $\sim$90~days (see Figs~\ref{fig:res} \&~\ref{fig:lifet}). The behavior of the active regions on the photosphere of Kepler-210 (i.e., the shift from small asterographic latitudes to higher and vice versa) is comparable to the migration of sunspots during an 11-year solar magnetic cycle (see Fig.~\ref{fig:ls-seasons}). We estimate that the duration of this phenomenon on Kepler-210 is similar, or somewhat longer, than the total {\it Kepler} observation time (i.e., $\sim$4 years), however, additional long-term observations are clearly needed to check whether this behavior of Kepler-210 as observed by {\it Kepler} is actually periodic and indeed the result of a magnetic cycle.
16
7
1607.08065
We use high accuracy photometric data obtained with the Kepler satellite to monitor the activity modulations of the Kepler-210 planet host star over a time span of more than four years. Following the phenomenology of the star's light curve in combination with a five spot model, we identify six different so-called spot seasons. A characteristic, which is common in the majority of the seasons, is the persistent appearance of spots in a specific range of longitudes on the stellar surface. The most prominent period of the observed activity modulations is different for each season and appears to evolve following a specific pattern, resembling the changes in the sunspot periods during the solar magnetic cycle. Under the hypothesis that the star exhibits solar-like differential rotation, we suggest differential rotation values of Kepler-210 that are similar to or smaller than that of the Sun. Finally, we estimate spot life times between ~60 days and ~90 days, taking into consideration the evolution of the total covered stellar surface computed from our model.
false
[ "spots", "differential rotation values", "the total covered stellar surface", "Sun", "days", "solar-like differential rotation", "six different so-called spot seasons", "the stellar surface", "longitudes", "high accuracy photometric data", "a five spot model", "the Kepler-210 planet host star", "the solar magnetic cycle", "consideration", "our model", "the observed activity modulations", "a specific range", "a specific pattern", "combination", "Kepler" ]
7.891151
12.559591
-1
12516895
[ "Frazin, Richard A." ]
2016SPIE.9909E..3XF
[ "Empirical Green's function approach for utilizing millisecond focal and pupil plane telemetry in exoplanet imaging" ]
2
[ "Univ. of Michigan (United States)" ]
[ "2016arXiv160804616F", "2019PASP..131k4506W" ]
[ "astronomy", "physics" ]
4
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "2009ApJ...701..804H", "2010SPIE.7736E..1JM", "2012SPIE.8453E..0TF", "2013ApJ...767...21F", "2013ApJ...767..100C", "2014ApJ...792...97M", "2014SPIE.9145E..3QF", "2014SPIE.9148E..17F", "2015A&A...581A..80R", "2015A&A...583A.102C", "2015ApJ...799...87B", "2015ApJ...808..149S", "2015PASP..127..445B", "2016JOSAA..33..712F" ]
[ "10.1117/12.2230786", "10.48550/arXiv.1607.08586" ]
1607
1607.08586_arXiv.txt
Direct imaging of exoplanetary systems is difficult due to the high contrast in brightness between the planet and the star, resulting in a situation in which the planetary image is confounded with the transient details of the telescope's point spread function (PSF). To date, all methods for direct imaging of exoplanets have relied on subtracting the PSF from the observed image to obtain the final science image. Typically, the PSF subtraction is achieved through differential imaging methods, which have been applied extensively in the recent literature. The fundamental difficulties associated with differential imaging methods, most notably angular differential imaging (ADI) and spectral differential imaging (SDI), have been reviewed in [\citenum{Frazin13,Frazin14,Frazin16a,Frazin16b,Marois_SOSIE,Rameau_ADI_SDI_limits15}] and references therein. Briefly, SDI is nearly useless close to the host star because the difference in the point spread function (PSF) at two wavelengths is proportional to the distance from the center of the image, and ADI assumes that the optical aberrations are not changing in time. Since differential imaging methods require estimating the speckle background from the final image, they suffer from a severe statistical penalty close to center \cite{Mawet_SmallNumStatsSpeckle14}, which is the most scientifically fruitful part of the image \cite{Stark_ExoEarthYield14,Brown_PlanetSearch15}. It is the author's belief that the combination of millisecond pupil plane and focal plane telemetry will lead to a large improvement in contrast over differential imaging techniques, as it leverages a vastly larger and richer data set than methods that use standard exposure times, which average over the atmospheric turbulence. Such methods are becoming practical due to a new generation of ultra-low noise IR cameras capable of kHz readouts, such as the SWIR single photon detector, SAPHIRA eAPD and the MKIDS \cite{SWIR_detector14,SELEX_APD12,Saphira_eAPD14,Mazin_MKIDS14}. It is important to emphasize that the author's millisecond imaging techniques can be generalized to take advantage of essentially all constraints on problem proposed to date. These constraints include those imposed by diurnal rotation (used by ADI), multi-wavelength observations (used by SDI), as well polarization (used in polarization differential imaging \cite{Hinkley_PDI09}). The basic concept of the author's approach was explained in [\citenum{Frazin13,Frazin14}]. Essentially, these papers provide a system of equations that express the instantaneous image measured in the science camera (SC) in terms of: \begin{itemize} \item{the residual phase [which is measured by the wavefront sensor (WFS)]} \item{the planetary image} \item{non-common path aberrations (NCPA).} \end{itemize} \begin{wrapfigure}{r}{0pt} \includegraphics[width=.49\linewidth,clip=]{fig_optical_system.eps} \caption{\small Schematic diagram of an astronomical telescope with a closed-loop AO system and a coronagraph. Modified from [\citenum{Hinnen_H2control}].} \label{fig_schematic} \vspace{-4mm} \end{wrapfigure} \noindent This system of equations forms the basis of a regression problem in which the NCPA and the planetary image are estimated self-consistently. The information required to successfully perform the estimation is provided by the variability of the residual phase (so it is closely related to other phase diversity methods), although multi-wavelength and diurnal rotation constraints could be employed as well. In [\citenum{Frazin16a,Frazin16b}], the author extended this approach to include non-pupil plane NCPA, which requires summing the contributions from unknown aberrations in various optical planes. The resulting expression is very complicated and computationally demanding because it requires propagation operators between the various optical planes. Here, the problem is formulated in full rigor, but in a way that is far more computationally tractable and should be suitable for parallel computation. The disadvantage of this approach is that the model has more degrees of freedom, but the richness of the millisecond data streams may well afford such a luxury.
Standard differential imaging methods such as ADI and SDI have a variety of systematic errors and are particularly ineffective close to the star, where biggest returns on exoplanet science are expected.\cite{Stark_ExoEarthYield14,Brown_PlanetSearch15}. The region within about $3 \lambda/D$\ (where $D$\ is the telescope diameter) is especially challenging due to the large angular displacement required for ADI, creating much incentive to explore methods beyond differential imaging. Fortunately, millisecond focal plane telemetry is now becoming practical due to a new generation of near-IR detector arrays with sub-electron noise that are capable of kHz readout rates. Combining these data with those simultaneously available from the WFS allows the possibility of self-consistently determining the optical aberrations (the cause of quasi-static speckles) and the planetary image.\cite{Frazin13,Frazin16a,Frazin16b}. Explicitly solving for the various aberrations in the optical system requires accounting for aberrations in various optical planes, leading to rather difficult and expensive computations, as many Fresnel propagations may be required.\cite{Frazin16a,Frazin16b}. Here, an attractive alternative is presented. Instead of explicitly solving for the aberrations in a number of planes, one may instead assume that the diffraction problem can be solved by a Green's function with an unknown component that accounts for the aberration, including the non-common path component. This unknown part of the Green's function is ``empirical Green's function,'' or EGF. Here, it was shown that estimating the EGF leads to a straightforward set of regression equations that avoid the plane-by-plane propagations shown in [\citenum{Frazin16a,Frazin16b}]. Furthermore, the EGF regressions can be made highly parallel due to their block-diagonal structure, which should make them practical. It was also shown that the EGF can be generalized to treat unknown polarization effects by employing a similar ``empirical Green's Tensor'' (EGT). The disadvantage of the EGF approach compared to explicit estimation the aberrations in multiple planes is that the EGF has many more degrees of freedom, although the richness of the millisecond datastreams may well afford the ability to estimate the larger number of parameters. It may be that as the field develops, the EGF (or EGT) would be most useful as an exploratory technique used guide and/or supplement techniques that solve for the aberrations explicitly.
16
7
1607.08586
Millisecond focal plane telemetry is now becoming practical due to a new generation of near-IR detector arrays with sub-electron noise that are capable of kHz readout rates. Combining these data with those simultaneously available from the wavefront sensing system allows the possibility of self-consistently determining the optical aberrations (the cause of quasi-static speckles) and the planetary image. This approach may be especially advantageous for finding planets within about 3λ/D of the star where differential imaging is ineffective. As shown in a recent article by the author (J. Opt. Soc. Am. A., 33, 712, 2016), one must account for unknown aberrations in several non-conjugate planes of the optical system, which, in turn, requires ability to computational propagate the field between these planes. These computations are likely to be difficult to implement and expensive. Here, a far more convenient alternative based on empirical Green's functions is provided. It is shown that the empirical Green's function (EGF), which accounts for all multi-planar, non-common path aberrations, and results in a much more tractable and highly parallel computational problem. It is also shown that the EGF can be generalized to treat polarization, resulting in the empirical Green's tensor (EGT).
false
[ "several non-conjugate planes", "sub-electron noise", "quasi-static speckles", "kHz readout rates", "Millisecond focal plane telemetry", "all multi-planar, non-common path aberrations", "unknown aberrations", "differential imaging", "J. Opt", "turn", "ability", "Green", "EGT", "the optical aberrations", "results", "the optical system", "the planetary image", "empirical Greens functions", "these planes", "the wavefront sensing system" ]
14.307842
10.659005
10
1444666
[ "Hagen, Lea M. Z.", "Seibert, Mark", "Hagen, Alex", "Nyland, Kristina", "Neill, James D.", "Treyer, Marie", "Young, Lisa M.", "Rich, Jeffrey A.", "Madore, Barry F." ]
2016ApJ...826..210H
[ "On the Classification of UGC 1382 as a Giant Low Surface Brightness Galaxy" ]
35
[ "Department of Astronomy and Astrophysics, The Pennsylvania State University, University Park, PA 16802, USA ; Institute for Gravitation and the Cosmos, The Pennsylvania State University, University Park, PA 16802, USA", "Observatories of the Carnegie Institution for Science, 813 Santa Barbara Street, Pasadena, CA 91101, USA", "Department of Astronomy and Astrophysics, The Pennsylvania State University, University Park, PA 16802, USA ; Institute for Gravitation and the Cosmos, The Pennsylvania State University, University Park, PA 16802, USA", "Physics Department, New Mexico Institute of Mining and Technology, Socorro, NM 87801, USA ; Netherlands Institute for Radio Astronomy (ASTRON), Postbus 2, NL-7990 AA Dwingeloo, the Netherlands ; National Radio Astronomy Observatory, Charlottesville, VA 22903, USA", "California Institute of Technology, Pasadena, CA 91125, USA", "Aix Marseille Université, CNRS, Laboratoire d'Astrophysique de Marseille, UMR 7326, 38 rue F. Joliot-Curie, F-13388 Marseille, France", "Physics Department, New Mexico Institute of Mining and Technology, Socorro, NM 87801, USA", "Observatories of the Carnegie Institution for Science, 813 Santa Barbara Street, Pasadena, CA 91101, USA; Infrared Analysis and Processing Center, California Institute of Technology, Pasadena, CA 91125, USA", "Observatories of the Carnegie Institution for Science, 813 Santa Barbara Street, Pasadena, CA 91101, USA" ]
[ "2016A&A...593A.126B", "2017ASSL..430..249S", "2017IAUS..325..379M", "2018MNRAS.480L..18Z", "2018MNRAS.481.3534S", "2019MNRAS.485..796M", "2019MNRAS.489.4669S", "2019arXiv190601657R", "2019sf2a.conf..249J", "2020A&A...637A..21J", "2020A&A...640A..38R", "2020ApJ...894..119D", "2020ApJ...900..142W", "2020MNRAS.496.3996K", "2020arXiv200414458S", "2020arXiv200613956L", "2021ApJ...914..104B", "2021ApJ...923...65F", "2021ApJ...923..273P", "2021MNRAS.502.5370W", "2021MNRAS.503..830S", "2022MNRAS.509.3148W", "2022aems.conf..395S", "2022arXiv220107588M", "2023A&A...676A..41J", "2023AJ....165..197G", "2023ApJ...959..105D", "2023MNRAS.520L..85S", "2023MNRAS.523.1538K", "2023MNRAS.523.3991Z", "2023MNRAS.525.3016M", "2023arXiv231203601L", "2024A&A...681A.100J", "2024ApJ...963...86L", "2024RAA....24a5018Z" ]
[ "astronomy" ]
17
[ "galaxies: individual: UGC 1382", "Astrophysics - Astrophysics of Galaxies" ]
[ "1955ApJ...121..161S", "1963BAAA....6...41S", "1965TrAlm...5...87E", "1973ApJ...186..467O", "1974ITAC...19..716A", "1981ApJ...246..666T", "1983ApJS...52...89H", "1983ApJS...53..105R", "1984A&A...132..389P", "1985ApJS...59..115K", "1987A&A...173...59V", "1987AJ.....94...23B", "1987Sci...235.1367S", "1990ApJ...360..427B", "1990Sci...250..539U", "1992ApJ...388L..13H", "1992StaSc...7..457G", "1994A&AS..106..451D", "1994A&AS..108..491L", "1994AJ....107..530M", "1995AJ....109..558S", "1995ApJ...441...18M", "1997AJ....114.1858P", "1997ARA&A..35..267I", "1997ApJ...490..493N", "1998ApJ...500..525S", "1999ApJS..121..287H", "1999MNRAS.310..540A", "2000ApJ...532..308B", "2000ApJ...533..682C", "2001A&A...365....1M", "2001AJ....121.2358B", "2001AJ....121.2420S", "2001MNRAS.328..353N", "2002AJ....124..266P", "2002PhRvD..66j3511L", "2003AJ....126.1090C", "2003ApJS..149..289B", "2003MNRAS.344.1000B", "2003PASP..115..763C", "2004A&A...422L...5R", "2004MNRAS.347..277M", "2004MNRAS.350.1195M", "2005AJ....130..873J", "2005ApJ...619L...1M", "2005MNRAS.361...34D", "2006A&A...457..841I", "2006A&A...460..381C", "2006ApJ...650L..33P", "2006PASA...23..165M", "2006PASP..118.1711W", "2007A&A...468...33E", "2007AJ....133.1085B", "2007ApJ...664..204S", "2007ApJS..173..267S", "2007ApJS..173..293W", "2007ApJS..173..538T", "2008ApJ...681..244B", "2008MNRAS.383.1223M", "2009A&A...504..807R", "2009ARA&A..47..159B", "2009ApJS..182..543A", "2009MNRAS.394..340G", "2010A&A...516A..11L", "2010AJ....139..315W", "2010AJ....139.2097P", "2010AJ....140.1194B", "2010AJ....140.1868W", "2010ApJ...714L.171T", "2010ApJ...714L.290S", "2010ApJS..186..427N", "2010MNRAS.404.1733S", "2010MNRAS.406L..90R", "2011A&A...532A..74B", "2011ApJ...728...74G", "2011ApJ...735..125S", "2011ApJ...737...47A", "2011ApJS..193...29A", "2011arXiv1102.0550B", "2012ApJ...745...34M", "2012ApJS..199...26H", "2013JApA...34...19D", "2013MNRAS.433.2986W", "2014MNRAS.437.3072K", "2015ApJ...801...97S", "2015ApJS..219...12A", "2015MNRAS.446.1512H" ]
[ "10.3847/0004-637X/826/2/210", "10.48550/arXiv.1607.02147" ]
1607
1607.02147_arXiv.txt
\label{sec-intro} Giant low surface brightness (GLSB) galaxies are the most extreme low surface brightness (LSB) disk galaxies and are the largest isolated galaxies known to exist. Although massive $(L \sim L^*)$ and gas rich ($\text{M}_\text{gas} > 10^{10}$~\msun), because they have disk scale lengths in excess of 10~kpc, they also have low gas surface densities and star formation efficiencies \citep{sprayberry95, impey97, matthews01}. Their rotation curves flatten near $V_\text{max} \sim$300~km/s and are dark matter (DM) dominated with DM fractions $> 0.7$ \citep{lelli10, buta11}. Despite the enormous size and luminosity of GLSB galaxies, their diffuse nature makes them difficult to detect and are assumed to be highly underrepresented in catalogs \citep{impey97}. Their contribution to the luminosity density of the universe remains unclear. Their origins have implications for the success of $\Lambda$CDM and hierarchical formation at low densities. GLSB galaxies are not simple `pure' low surface brightness systems. Rather, a defining characteristic is that they have both a normal high surface brightness (HSB) central component (typically an early type disk) which is embedded in a massive extended diffuse disk component \citep{lelli10, sprayberry95, barth07}. Because star formation is usually present in the extended disks of GLSB galaxies \citep{boissier08}, they can be considered larger versions of the more recently defined category of Type~1 extended ultraviolet (XUV) disk galaxies \citep{thilker07}, in which UV emission is seen at distances well beyond the classical star formation threshold surface density. GLSB galaxies, but for their large scale, are also similar to the population of low-mass early type galaxies (i.e., elliptical and lenticular galaxies) that show low levels of star formation in the outer regions, which may be the result of recent accretion of lower mass galaxies \citep{moffett12, salim10}. There is no definitive formation scenario for GLSB galaxies, but most agree that a low density environment is required in order to build and keep such enormous, organized, tenuous, and seemingly undisturbed extended disks. Although often described as simply unevolved gas-rich disks due to their low star formation efficiency \citep{bothun87, hoffman92}, the dual HSB inner region and LSB extended disk suggest a more complicated history which may involve both a rapid disk formation and a late collapse of a low-amplitude density perturbation \citep{impey97} or the tidal disruption of dwarf galaxies \citep[e.g.][]{penarrubia06}. There are likely several mechanisms at work simultaneously. However they form, the relative isolation and low star formation efficiency suggests that they are not evolving rapidly at present. The prototypical GLSB galaxy, Malin~1, discovered by \citet{bothun87}, has an extrapolated disk central surface brightness of $\mu_R(0) =$24.7 mag/arcsec$^2$, a staggering disk scale length of 57~kpc \citep[for $h = 70$;][]{moore06} and an absolute magnitude of $M_V = -22.9$ \citep{pickering97}. Malin~1's HI disk has a mass of $10^{11} M_\odot$ and extends to a radius of 110~kpc \citep{lelli10, pickering97}. Using Hubble Space Telescope imaging, \citet{barth07} confirmed that the inner 10~kpc of Malin~1 hosts a SB0/a disk of normal size and surface brightness. \citet{boissier08} has classified Malin~1 as having a Type~1 XUV disk. Although more than a dozen systems are now considered to be GLSB galaxies \citep{matthews01, bothun90, sprayberry95}, no other system has been reported with properties as extreme as the prototypical Malin~1. In this article, we describe UGC~1382, which is nearly identical in terms of scale and other physical properties to Malin~1. However, at less than 1/4 the distance to Malin~1, it is significantly closer. This allows a detailed multi-wavelength investigation of a true Malin~1-like GLSB galaxy at much smaller spatial scales, with the goal of constraining the formation and evolution of these extreme systems. \ugc\ has been classified as an elliptical in many optical surveys \citep{tonry81, laurikainen94, huchra99, doyle05, sanchez11, huchra12}. Several surveys looking for morphological features, such as stellar rings and bars, did not detect anything other than a simple bulge-dominated galaxy \citep{meyer04, nair10, baillard11}. It has spectroscopically measured recession velocities ranging between 5550 and 5770~km/s \citep{huchra83, huchra99, meyer04, garcia09, aihara11}. We find a 21~cm systemic radial velocity of 5591~km/s (see \S\ref{sec-HI}) and adopt a distance of 80~Mpc \citep{wright06} in this paper; this gives a scale of about 380~pc/arcsec or 23~kpc/arcmin. UGC~1382 may be in a small group; there are three known galaxies within 1.5~Mpc. UGC~1382 was found to have $5 \times 10^9$~\msun\ of HI gas \citep{garcia09}, which is approximately 13\% of the stellar mass \citep{west10}. The only hint that it may be more noteworthy was the suggestion of an extended HI disk \citep{garcia09}, though no analysis of such a disk was undertaken. UGC~1382 came to our attention during an investigation of star formation in early type galaxies. We noticed that it contained a set of very extended spiral arms in ultraviolet (UV) imaging from the Galaxy Evolution Explorer \citep[\galex;][]{martin05}. Further investigation revealed that this system is not an elliptical galaxy, but is in fact a GLSB galaxy composed of a HSB lenticular core and an 80~kpc radius LSB disk. In order to better understand this unusual galaxy, we have assembled a set of multiwavelength data, ranging from radio to far-ultraviolet, which we present in Section~\ref{sec-data}. In Section~\ref{sec-morph}, we discuss the galaxy's morphology, surface brightness profiles, HI gas content, star formation efficiency, LSB characteristics, and environment. We derive the dark matter content of the galaxy in Section~\ref{sec-dm}. We then use the multiwavelength photometric data to model the spectral energy distribution (SED) of the galaxy, its HSB lenticular component, and its extended LSB disk in Section~\ref{sec-model}. In Section~\ref{sec-green}, we examine the past and future evolution of UGC~1382 based on both its morphology and modeled physical parameters. We present possible formation scenarios in Section~\ref{sec-formation}. Finally, we summarize our results in Section~\ref{sec-summary}. We use flat $\Lambda$CDM cosmology with $\Omega_\Lambda = 0.7$, $\Omega_\text{M} = 0.3$, and $H_0=70$~km/s/Mpc throughout.
16
7
1607.02147
We provide evidence that UGC 1382, long believed to be a passive elliptical galaxy, is actually a giant low surface brightness (GLSB) galaxy that rivals the archetypical GLSB Malin 1 in size. Like other GLSB galaxies, it has two components: a high surface brightness disk galaxy surrounded by an extended low surface brightness (LSB) disk. For UGC 1382, the central component is a lenticular system with an effective radius of 6 kpc. Beyond this, the LSB disk has an effective radius of ∼38 kpc and an extrapolated central surface brightness of ∼26 mag arcsec<SUP>-2</SUP>. Both components have a combined stellar mass of ∼8 × 10<SUP>10</SUP> M <SUB>⊙</SUB>, and are embedded in a massive (10<SUP>10</SUP> M <SUB>⊙</SUB>) low-density (&lt;3 M <SUB>⊙</SUB> pc<SUP>-2</SUP>) HI disk with a radius of 110 kpc, making this one of the largest isolated disk galaxies known. The system resides in a massive dark matter halo of at least 2 × 10<SUP>12</SUP> M <SUB>⊙</SUB>. Although possibly part of a small group, its low-density environment likely plays a role in the formation and retention of the giant LSB and HI disks. We model the spectral energy distributions and find that the LSB disk is likely older than the lenticular component. UGC 1382 has UV-optical colors typical of galaxies transitioning through the green valley. Within the LSB disk are spiral arms forming stars at extremely low efficiencies. The gas depletion timescale of ∼10<SUP>11</SUP> years suggests that UGC 1382 may be a very-long-term resident of the green valley. We find that the formation and evolution of the LSB disk in UGC 1382 is best explained by the accretion of gas-rich LSB dwarf galaxies.
false
[ "HI disk", "galaxies", "a high surface brightness disk galaxy", "<", "other GLSB galaxies", "LSB", "a giant low surface brightness", "∼8 ×", "an extended low surface brightness", "the largest isolated disk galaxies", "∼38 kpc", "gas-rich LSB dwarf galaxies", "size", "∼26 mag arcsec", "lt;3 M <SUB>⊙</SUB", "an extrapolated central surface brightness", "spiral arms forming stars", "the LSB disk", "UGC", "the giant LSB and HI disks" ]
11.199212
7.354337
194
12500425
[ "Sasaki, Toru", "Matsushita, Kyoko", "Sato, Kosuke", "Okabe, Nobuhiro" ]
2016PASJ...68...85S
[ "X-ray observations of a subhalo associated with the NGC 4839 group infalling toward the Coma cluster" ]
9
[ "Department of Physics, Tokyo University of Science, 1-3 Kagurazaka, Shinjuku-ku, Tokyo 162-8601, Japan", "Department of Physics, Tokyo University of Science, 1-3 Kagurazaka, Shinjuku-ku, Tokyo 162-8601, Japan", "Department of Physics, Tokyo University of Science, 1-3 Kagurazaka, Shinjuku-ku, Tokyo 162-8601, Japan", "Department of Physical Science, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526, Japan; Hiroshima Astrophysical Science Center, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526, Japan" ]
[ "2017A&A...605A..25E", "2018AJ....156..224J", "2018ApJ...866...48U", "2019ApJ...874..112S", "2019MNRAS.485.2922L", "2019SSRv..215....7W", "2023ApJ...944...51O", "2024ApJ...966..236L", "2024RAA....24d5016F" ]
[ "astronomy" ]
7
[ "galaxies: clusters: individuals (Coma cluster", "NGC 4839 group)", "galaxies: clusters: intracluster medium", "X-rays: galaxies: clusters", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
[ "1972ApJ...176....1G", "1981A&A...100..323B", "1982MNRAS.198.1007N", "1992A&A...259L..31B", "1995MNRAS.275..720N", "1998SSRv...85..161G", "1999ApJS..125...35S", "2000ApJ...541..542M", "2001A&A...365L..74N", "2001ApJ...556L..91S", "2002ApJ...567L..27M", "2002PASJ...54..327K", "2003A&A...400..811N", "2003ApJ...591.1220L", "2004MNRAS.348..811T", "2004MNRAS.350.1397T", "2006ApJ...640..691V", "2006ApJ...648L.109C", "2006ApJ...650..102A", "2007PASJ...59S.113I", "2007PhR...443....1M", "2008A&A...478..575K", "2008A&A...486..359L", "2008PASJ...60..345O", "2008SSRv..134...93F", "2009ApJ...693.1142S", "2010A&A...516A..32G", "2010ApJ...713..291O", "2010ApJ...715.1143F", "2010MNRAS.406.1721R", "2011ApJ...741..116O", "2013ApJ...775....4S", "2013MNRAS.433.1701O", "2013PASJ...65...16A", "2013PASJ...65...89A", "2014A&A...570A.119E", "2014ApJ...784...90O", "2014MNRAS.444..629R", "2015A&A...582A..87A", "2015ApJ...806..123S", "2015MNRAS.448.2971I", "2015PASJ...67..113I", "2016ApJ...817...24M" ]
[ "10.1093/pasj/psw078", "10.48550/arXiv.1607.07554" ]
1607
1607.07554_arXiv.txt
Galaxy clusters have formed through mergers of clusters and/or accretions of smaller groups and galaxies. X-ray observations have revealed features such as shocks, cold fronts, and gas stripping tails, around accreting galaxy groups to clusters (e.g., \cite{Markevitch2000,Markevitch2002, Markevitch2007,Russell2010,Ghizzardi2010,Russell2014, Eckert2014, Ichinohe2015} ). Shock fronts are expected to occur in front of a merging subhalo and to heat the intracluster medium (ICM). Radio relics, which are diffuse synchrotron radio emissions, are thought to be tracers of shock structures \citep{Ferrari2008} and recent {\it XMM-Newton}, {\it Chandra} and {\it Suzaku} observations have found jumps in temperature across these radio relics (e.g., \cite{Finoguenov2010, Akamatsu2013a, Itahana2015, Akamatsu2015}). \citet{Markevitch2000} found an edge \textcolor{blue}{in} surface brightness where the temperature jumped, and the gas pressure across the edge was continuous. Such a feature is called a `cold front', and is caused by gas sloshing when a minor merger oscillates the cool dense core of a cluster (e.g., \cite{Ascasibar2006}). In addition, the accreting galaxies and groups experience the stripping of gas by ram pressure (e.g., \cite{Gunn1972}). The gas stripped from the accreting objects looks like a tail structure in the X-ray band. Weak-lensing measurement is a direct, powerful technique to probe dark matter distribution of clusters without resting any assumptions of dynamical states. Since the lifetime of the dark matter subhalos is longer than that of gas subhalos \citep{Tormen2004}, weak-lensing enables us to measure the dark matter mass of the subhalos whose interaction triggers the gas features discussed above. Therefore, a joint weak-lensing and X--ray study is powerful to understand cluster mergers. To date, several joint studies of major mergers have revealed the interaction between ICM and the dark matter subhalo \citep{Clowe2006, Okabe2008, Okabe2011, Okabe2015, Medezinski2015}. However, joint studies on minor mergers are not yet sufficient. A recent weak-lensing study (\cite{Okabe2014}; see also \cite{Okabe2010}) for the Coma cluster detected 32 small subhalos whose mass was greater than $\sim 3\times10^{12}~h^{-1} M\solar$. The ratio between the subhalo mass and the virial mass reaches down to a few $10^{-3}$. Detections of such small subhalos was achievable by the large apparent size, because the cluster is very close to us. There are four massive subhalos whose masses are greater than $\sim 1\times10^{13}~M\solar$, and labeled ``ID~1'', ``ID~2'', ``ID~9'', and ``ID~32'' in \citet{Okabe2014}. The massive subhalo signals excluding the ``ID~2'' in the Coma cluster, as well as stacked signals from all subhalos, were well represented with a truncated Navarro-Frenk-White (NFW) mass model \citep{Navarro1995} as expected from a tidal destruction model \citep{Okabe2014}. These subhalos excluding the ``ID~9'' were observed with Suzaku \citep{Sasaki2015}. While excess X-ray emission from the ``ID~1'' and ``ID~32'' subhalos was detected, the ``ID~2'' did not have excess X-ray emission. The gas mass to weak-lensing mass ratios of the ``ID~1'' and ``ID~32'' subhalos were 0.001, which is one to two orders of magnitude smaller than those of regular groups. This supports that ram pressure with typical infall velocities can strip significant amounts of gas from the infalling subhalos. In this paper, we describe the thermodynamics around the third massive subhalo "ID~9" whose mass and truncation radius ($r_t$) are $(1.58\pm0.26)\times10^{13}~M\solar$ and 3.43\arcmin$_{-0.47}^{+0.28}$, respectively, as detected by weak-lensing analysis \citep{Okabe2014}. The subhalo ``ID~9" is associated with the NGC~4839 group, which is the most famous merging substructure in the Coma cluster. It is located at a projected distance of approximately 50$\arcmin$ ($\sim r_{500}$ \footnote{$r_{500}$ is the radius within which the mean density of the cluster is 500 times the critical density of the universe.}). The NGC 4839 group has an X-ray tail toward the southwest or opposite to the direction of the Coma cluster center \citep{Briel1992}. With {\it XMM-Newton}, \citet{Neumann2001} reported complex temperature structures around the NGC~4839 group, with indications of a bow shock and of ram pressure stripping. At the southwest of the NGC~4839 group, a radio relic was discovered around the virial radius of the Coma cluster \citep{Ballarati1981}. The X-ray observations found that there is temperature discontinuity at the relic, which corresponds to a shock with the {\cal M}ach number $\sim$ 2 \citep{Ogrean2013, Akamatsu2013b}. We summarize the observation logs and data reductions in section \ref{sec:obs}. Section \ref{sec:ana} shows an X-ray image around subhalo ``ID~9" with the mass distribution derived from weak-lensing analysis, and results of spectral analysis. We discuss thermodynamics around the massive subhalo in section \ref{sec:dis}. In this paper, we use $\Omega_{m,0}=0.27, \Omega_{\Lambda}=0.73$, and $H_{0}=$70~km~s$^{-1}$~Mpc$^{-1}$. The redshift of the Coma cluster is $z$ = 0.0231 \citep{Struble1999}. At this redshift, $1\arcmin$ corresponds to 28.9~kpc. The solar abundance table is given by \citet{Lodders2003}. Unless noted otherwise, the errors are in the 68\% confidence region for the single parameter of interest.
We observed the third massive subhalo ``ID~9" detected from weak-lensing analysis \citep{Okabe2014} with {\it Suzaku}. By comparing the subhalo mass distribution with thermodynamics obtained by X-ray observation, we first investigated cluster evolution through the accretion of objects. While the X-ray peak is shifted approximately $1\arcmin$ away from the mass center, the NGC~4839 galaxy coincided with the X-ray peak. The X-ray tail was elongated toward the southwest direction or the outskirts of the Coma cluster, and the length was approximately $6~r_t$, $\sim$ 600~kpc. The temperatures of the central core and X-ray tail regions were approximately 5~keV. At a distance $r_t$ toward the north, corresponding to the head of the subhalo, we found a temperature jump, from 5 keV to 8--10 keV. While the temperature of the southern part was also 7~keV at 1-2~$r_t$, temperature profile beyond 2~$r_t$ decreased to 4 keV. At $2~r_t$ in the southern part, we estimated a Mach number to be ${\cal M}=1.73^{+0.25}_{-0.16}$ under the Rankine-Hugoniot condition. The gas mass fraction of the subhalo within $r_t$ was approximately 0.7\%. It is approximately 5 times smaller than that of regular groups and poor clusters. With the infall velocity estimated from the weak-lensing mass of the Coma cluster \citep{Okabe2010}, the ram pressure was comparable with the gravitational force per unit area. Assuming the Kelvin-Helmholtz instabilities, the removed gas mass was estimated to be approximately $3\times10^{11}~M\solar$. Assuming that the removed gas had been within $r_t$ before infalling to the Coma cluster, the gas mass fraction of the subhalo had been comparable with those of regular groups. \bigskip We thank the anonymous referee for careful reading the manuscript and providing valuable comments. We gratefully acknowledge all members of the {\it Suzaku} operation, calibration, and software development teams. TS is supported by JSPS Research Fellows. We acknowledge the support of Grants-in-Aid for Scientific Research from the MEXT, No. 25400235(K. M.), 25800112 (K. S.), and 26800097(N. O.). This work was supported by the Funds for the Development of Human Resources in Science and Technology under MEXT, Japan, and Core Research for Energetic Universe in Hiroshima University (the MEXT program for promoting the enhancement of research universities, Japan). \appendix
16
7
1607.07554
We report on Suzaku X-ray observations of the dark subhalo associated with the merging group of NGC 4839 in the Coma cluster. The X-ray image exhibits an elongated tail toward the southwest. The X-ray peak shifts approximately 1' away from the weak-lensing mass center toward the opposite direction of the Coma cluster center. We investigated the temperature, normalization, pressure, and entropy distributions around the subhalo. Excluding the X-ray tail, the temperature beyond the truncation radius is 8-10 keV, which is twice as high as that of the subhalo and the X-ray tail. The pressure is nearly uniform, excluding the southern part of the subhalo at two times of the truncation radius. We computed the gas mass within the truncation radius and the X-ray tail. While the gas fraction within the truncation radius is about five times smaller than that of regular groups, the gas mass in the subhalo and the X-ray tail to weak-lensing mass ratio is consistent with that of regular groups. Assuming an infall velocity of 2000 km s<SUP>-1</SUP>, the ram pressure is 1.4 times greater than the gravitational force per unit area. Assuming the Kelvin-Helmholtz instabilities, the total lost mass is approximately 3 × 10<SUP>11</SUP> M<SUB>⊙</SUB>. If this gas had originally been within the truncation radius, the gas mass fraction of the subhalo would have been comparable with those of regular groups before infalling to the Coma cluster.
false
[ "Coma", "the Coma cluster center", "regular groups", "the Coma cluster", "Suzaku X-ray observations", "unit area", "the opposite direction", "the X-ray tail", "The X-ray peak shifts", "The X-ray image", "the weak-lensing mass center", "entropy distributions", "the gas mass fraction", "pressure", "NGC", "the truncation radius", "the merging group", "the total lost mass", "the gas mass", "an elongated tail" ]
13.140255
4.814026
163
12472995
[ "Şahin, T.", "Lambert, David L.", "Klochkova, Valentina G.", "Panchuk, Vladimir E." ]
2016MNRAS.461.4071S
[ "HD 179821 (V1427 Aql, IRAS 19114+0002) - a massive post-red supergiant star?" ]
16
[ "Faculty of Science, Space Sciences and Technologies Department, Akdeniz University, 07058 Antalya, Turkey; The W. J. McDonald Observatory, Department of Astronomy, The University of Texas, Austin, TX 78712, USA", "The W. J. McDonald Observatory, Department of Astronomy, The University of Texas, Austin, TX 78712, USA", "Special Astrophysical Observatory, Nizhnij Arkhyz, Karachai-Cherkessia 369167, Russia", "Special Astrophysical Observatory, Nizhnij Arkhyz, Karachai-Cherkessia 369167, Russia; National Research University ITMO, St Petersburg 197101, Russia" ]
[ "2017ASPC..510..121K", "2018A&A...613A..33C", "2018ARep...62...19K", "2018AstBu..73...52K", "2018AstL...44..457I", "2019A&A...631A..48V", "2019Ap&SS.364...18P", "2019AstBu..74...41K", "2019AstBu..74..475K", "2020ApJ...899...41S", "2022MNRAS.511.4360K", "2023Galax..11...76K", "2023PARep...1...54S", "2024A&A...686A.270K", "2024arXiv240505309K", "2024arXiv240607196G" ]
[ "astronomy" ]
13
[ "stars: abundances", "stars: late-type", "stars: massive", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1973PhDT.......180S", "1975ApJ...195..137S", "1977MNRAS.180..631S", "1979ApJ...232..409H", "1980A&A....84..361B", "1982ApJ...261..736H", "1983ApJ...272...99W", "1984sfat.book.....K", "1985Ap.....22..203B", "1985QJRAS..26....1I", "1986ApJ...311..345Z", "1986MNRAS.222..809H", "1987A&A...173L..11L", "1987SvAL...13..214D", "1988Ap.....28..197B", "1988Ap.....28..202B", "1988LNP...305...94B", "1989A&AS...77..137T", "1992oasp.book.....G", "1993A&A...269..231V", "1993A&A...275...67B", "1994ApJS...94..221N", "1995ApJ...449..839V", "1995PASP..107..251T", "1996MNRAS.282.1171Z", "1996atpc.book.....W", "1997A&A...318..269M", "1998AJ....115.1640M", "1998PASJ...50..629T", "1998SPIE.3355..387T", "1999AJ....117.1834R", "1999ApJ...521..753T", "1999ApJ...525L.113J", "2000A&A...359..138T", "2001ApJ...556..408J", "2001ApJ...556..452L", "2001AstL...27..156A", "2002A&A...389..519A", "2003A&A...400..613L", "2003ApJS..148..543D", "2003IAUS..210P.A20C", "2003RMxAA..39....3A", "2004PASP..116..693M", "2005ApJ...622.1058D", "2005PASJ...57..751T", "2006A&A...448..641D", "2006A&A...456.1181L", "2006ApJ...652.1626F", "2006JPCRD..35.1669F", "2007A&A...465..457C", "2007A&A...466..327B", "2007AJ....133.2464L", "2007ApJ...667.1267S", "2007AstBu..62..105A", "2007uasb.conf..179P", "2008ApJ...680.1256H", "2008AstBu..63..386K", "2008BaltA..17...87K", "2008JPCRD..37.1285K", "2008PhDT........37S", "2009ARA&A..47..481A", "2009ASPC..412...17O", "2009AstL...35..764A", "2009JPCRD..38..171P", "2011A&A...528A.103L", "2011AJ....142..136L", "2011MNRAS.410..612S", "2012MNRAS.423.3268K", "2012MNRAS.427...50L", "2012MNRAS.427..127B", "2013A&A...557A..26M", "2013ApJS..208...27W", "2014A&A...562A..71B", "2014AJ....147..137L", "2014dapb.book..169B", "2015A&A...573A..26S", "2015A&A...581A..70K", "2016A&A...586A.120O" ]
[ "10.1093/mnras/stw1586", "10.48550/arXiv.1607.00142" ]
1607
1607.00142_arXiv.txt
The spectral class of F-G luminous giants may encompass stars on two different evolutionary paths. Some stars may be massive stars evolving from the main sequence and some of these massive stars may now be in a post-red supergiant phase. Alternatively, other stars may be departing the asymptotic giant branch (AGB) evolving at roughly constant luminosity to hotter temperatures and the tip of the white dwarf cooling track. Unambiguous assignment of a F-G supergiant to the proper evolutionary path is not always immediately possible, even when a wide variety of observational techniques are applied and the electromagnetic spectrum is well sampled. HD\,179821, also known as V1427 Aql and IRAS\,19114+0002, remains a supergiant of uncertain heritage despite a lengthy literature and frequent investigations into its status. Advocates for a post-AGB origin include Za\v{c}s et al. (1996) and Reddy \& Hrivnak (1999) who gave weight to their measurements of overabundances of $s$-process nuclides. Others have stressed the star's distance as implied by its radial velocity and characteristics of its circumstellar gas and dust shell in suggesting that the star is a massive supergiant: see, for example, Jura \& Werner (1999) and Jura, Velusamy, \& Werner (2001) with the latter paper carrying the provocative title `What next for the likely presupernova HD 179821?'. Oudmaijer et al. (2009) confidently place HD\,179821 among massive stars in a post-red supergiant phase. An oft stated assertion is that a star's chemical composition provides clues to its evolutionary history. Certainly, one anticipates readily observable distinctive signatures between an evolved massive star and a mature post-AGB star (i.e., a star that left the AGB after experiencing many thermal pulses and extensive third dredge-up episodes). The massive star will have experienced mixing between envelope and interior at a minimum and, thus, readjustment of its surface C, N and possibly O abundances: a decrease in C abundance with an offsetting increase in N abundance is assured but with very few exceptions (Na, possibly) all other elements will retain their natal abundances. On the other hand, a mature post-AGB star will be markedly $s$-process enriched with a likely enrichment of C accompanying the $s$-process enrichment. The contrasting compositions of massive and post-AGB stars surely represent a testable proposition. Unfortunately, some stars appear to evolve off the AGB before the third dredge-up has altered the surface abundances of the $s$-process nuclides (see De Ruyter et al. 2006, and references therein). The paper is organized as follows: Section 2 discusses the high-resolution optical spectra obtained at the two observatories from 2008--2013; General properties of the spectra are discussed briefly in Section 3; Section 4 describes our abundance analysis and reanalyses previous analyses; Section 5 in a collection of concluding remarks places HD\,179821 in its likely evolutionary status as a massive star evolving to become a red supergiant and finally a Type II supernova.
There is no question but that HD\,179821 is a luminous star. Possible identifications of the star include two possibilities: (i) a massive star evolving from the main sequence at approximately constant luminosity to the red supergiant phase or on a post-red supergiant loop back to the blue or (ii) a lower mass star evolving at roughly constant luminosity from the AGB to the tip of the white dwarf cooling track. The mass-luminosity relations for these alternative identifications overlap for a range of masses. Above a certain luminosity, the more likely identification is a high mass star. This critical luminosity is set by the maximum mass -- the Chandrasekhar mass -- of a post-AGB star which may become a white dwarf. At the Chandrasekhar mass, the luminosity of a post-AGB star is M$_{bol} \simeq $-7.1 or $\log L/\Lsolar \simeq $4.7 (Wood, Bessell, \& Fox 1983). Therefore, if it can be shown that HD 179821's luminosity exceeds the latter limit by a clear margin, one may identify the star as a massive star. Estimations of HD\,179821's absolute magnitude from an apparent magnitude are fraught with uncertainty owing to an uncertain correction for interstellar extinction with the possibility of an additional correction for circumstellar extinction. As previous investigators of the star have appreciated (Reddy \& Hrivnak 1999; Kipper 2008; Oudmaijer et al. 2009), the absolute magnitude of HD 179821 may be estimated from the equivalent width (EW) of the O\,{\sc i} triplet at 7770-7774 \AA. The $M_V$ -- $EW$ calibration for warm supergiants comes from Kovtyukh, Gorlova \& Belik (2012) who assembled EW measurements for supergiants with known luminosities. Our measurement of the oxygen triplet's EW is 2.7 \AA. Kovtyukh et al.'s calibration has few stars with such strong EW but mild extrapolation of the calibration gives $M_V \simeq $-8.9 or $\log L/\Lsolar\simeq $5.5. This absolute magnitude is nearly two magnitudes brighter than the maximum for a post-AGB star and slightly fainter than the Humphreys-Davidson (Humphreys \& Davidson 1979) limit of M$_{\rm bol} \simeq $-9.5 for the most luminous warm Galactic supergiants such as $\rho$ Cas and HR 8752. Absolute luminosities for post-AGB stars are poorly known, in general. Arellano Ferro, Giridhar, Arellano (2003) incorporated five post-AGB stars into their calibration of $M_V$ -- $EW$ relation based primarily on normal supergiants. The two post-AGB star with estimated luminosities close to the luminosity limit for post-AGB stars had $EW$s of the triplet of 2.0 and 1.7 \AA, values similar to $EW$s of normal supergiants of the same luminosity. Thus, it appears that normal supergiants and post-AGB stars of type F-G satisfy similar $M_V$ -- $EW$ relations for the oxygen triplet. The effective temperature and surface gravity provide a check on the conclusion that HD 179821 is a massive star. By combining the relations $L \propto R^2T_{\rm eff}^4$ and $g \propto M/R^2$, one obtains \begin{equation} \hskip 1.5 cm log L/\Lsolar = log M/\Msolar + 4log T_{\rm eff} - log g - 10.61 \end{equation} On substituting T$_{\rm eff} = 7350$K, $\log g = 0.64$ and $\log L/\Lsolar = 5.5$, a mass $M = 19M_\odot$ is obtained. Stellar evolutionary tracks (e.g., Iben 1985) imply that such a luminosity is achieved at a higher mass -- say, 30$M_\odot$ -- but an adjustment of $\log g$ by just 0.2 dex provides just such a mass\footnote{An independent estimate of the mass was also provided by Parsec isochrones (Bressan et al. 2012) with the solar metal content of Z$_{\rm \odot} =$ 0.0152 and using a Z = 0.03 ([Fe/H] = log(Z/Z$_{\rm \odot}$)) isochrone (http://stev.oapd.inaf.it/cgi-bin/cmd-2.7). This isochrone, based on model atmosphere parameters reported in Section 4.2, provided a mass of 30 $\Msolar$ for HD\,179821.}. It remains to consider HD\,179821's composition in light of the proposed identification as a massive star which has evolved beyond core H-burning on the main sequence to a He-core burning warm supergiant (and possibly beyond this phase!). A massive star observed as a warm supergiant may be expected to have shuffled at a minimum its C and N abundances as CN-processed material reached the surface by rotationally-induced mixing and the first dredge-up with the latter a contributor if the star is now evolving to the blue after a red supergiant phase. Present C/N/O abundances are taken from our analysis summarized in Table 6 with the approximate non-LTE corrections listed in Section 4.6, i.e., C = 8.5, N= 9.2 and O = 8.9. On the assumption that initial abundances satisfied the condition [X/Fe] = 0.0 and [Fe/H] = +0.4, C= 8.9, N = 8.5 and O= 9.2 were the starting abundances; relative to these values N is clearly enriched, C and O are depleted. Assuming that the envelope has been mixed with CN-processed material from the interior, the CN-cycle's catalysts of C and N are required to be conserved. The initial (logarithmic) sum is 9.0 and the observed sum is 9.3 with conservation satisfied within the measurement uncertainties and including the rough corrections for non-LTE effects. For more severe mixing, ON-cycled products may be involved and the conserved quantity is the sum of the C, N and O abundances. In this case, the initial value is 9.4 and the observed value is 9.3 which surely represents a case of fortuitous agreement. (Note: the analysis does not consider enhancement of the surface He abundance during evolution.) Sodium overabundances in F-G supergiants have been the subject of several observational and theoretical investigations. As an observational reference point for HD 179821, we take the survey by Andrievsky et al. (2002) of Na abundances in F-G Ib-II giants which showed [Na/Fe] increasing with decreasing $\logg$ reaching [Na/Fe] $\simeq +0.3$ at $\logg = 1.0$. Non-LTE calculations show that the departures from LTE on the observed Na\,{\sc i} lines at 6154 \AA\ and 6160\AA\ are small ($\sim$ 0.1 dex) for supergiants with $\Teff \sim 7000$ \kelvin. The Na enrichment is attributed to operation of the H-burning Ne-Na chain in which $^{22}$Ne is converted to $^{23}$Na. Denissenkov (2005) argued that observed levels of Na enrichment required massive stars to undergo mixing between core and the radiative envelope in their main sequence progenitor. Inspection of Andrievsky et al.'s list of stars showed that just one was as Fe-rich as HD\,179821. A reference sample of four stars was selected with the conditions that $\Teff \geq 6500$ \kelvin\, and $\logg \le 1.0$. For this quartet, mean values are [Fe/H] $= -0.24$ and [Na/Fe] $= +0.37$ with a small star-to-star scatter. If we assume that the $^{22}$Ne abundance scales with [Fe/H], and the conversion of $^{22}$Ne to Na with subsequent mixing of Na to the surface are independent of metallicity, HD\,179821 with [Fe/H] $\simeq +0.4$ and [Na/H] also $\simeq +0.4$ for an unmixed star is expected to have [Na/Fe] $\sim 0.7$, a value close to the observed value in Table 6. An alternative identification of HD\,179821 as a post-AGB star might appear to be excluded on the grounds that key signatures of post-AGB stars are absent, i.e., a C-rich and $s$-process atmosphere and envelope are clearly not a feature of HD 179821. This exclusion supposes that all AGB stars experience third dredge-up (i.e., envelope enrichment with C and $s$-process nuclides) before evolving off the AGB. The case of RV Tauri stars as post-AGB stars provides common examples of without carbon and $s$-process enrichment. In such situations, when chemical composition is not a definitive way to distinguish massive evolved from post-AGB stars, other observed characteristics may be invoked. For example, the circumstellar CO expansion velocity for HD\,179821 exceeds the typical velocity for post-AGB stars and likely requires a very luminous (i.e., a massive) star in order to drive the expansion by radiation pressure (Jura et al. 2001; Oudmaijer et al. 2008). But, obviously, among the other characteristics the absolute luminosity plays a key role, one looks forward to a precise trigonometrical parallax. Until then, HD\,179821 may be considered to be a warm Galactic supergiant like $\rho$ Cas and HR 8752.
16
7
1607.00142
We have derived elemental abundances of a remarkable star, HD 179821, with unusual composition (e.g. [Na/Fe] = 1.0 ± 0.2 dex) and extra-ordinary spectral characteristics. Its metallicity at [Fe/H] = 0.4 dex places it among the most metal-rich stars yet analysed. The abundance analysis of this luminous star is based on high-resolution and high-quality (S/N ≈ 120-420) optical echelle spectra from McDonald Observatory and Special Astronomy Observatory. The data includes five years of observations over 21 epochs. Standard 1D local thermodynamic equilibrium analysis provides a fresh determination of the atmospheric parameters over all epochs: &lt;italic&gt;T<SUB>eff</SUB>&lt;/italic&gt; = 7350 ± 200 K, log g= +0.6 ± 0.3, and a microturbulent velocity ξ = 6.6 ± 1.6 km s<SUP>-1</SUP> and [Fe/H] = 0.4 ± 0.2, and a carbon abundance [C/Fe] = -0.19 ± 0.30. We find oxygen abundance [O/Fe] = -0.25 ± 0.28 and an enhancement of 0.9 dex in N. A supersonic macroturbulent velocity of 22.0 ± 2.0 km s<SUP>-1</SUP> is determined from both strong and weak Fe I and Fe II lines. Elemental abundances are obtained for 22 elements. HD 179821 is not enriched in s-process products. Eu is overabundant relative to the anticipated [X/Fe] ≈ 0.0. Some peculiarities of its optical spectrum (e.g. variability in the spectral line shapes) is noticed. This includes the line profile variations for H α line. Based on its estimated luminosity, effective temperature and surface gravity, HD 179821 is a massive star evolving to become a red supergiant and finally a Type II supernova.
false
[ "Special Astronomy Observatory", "McDonald Observatory", "log g= +0.6 ±", "Fe I and Fe II lines", "extra-ordinary spectral characteristics", "g= +0.6", "the spectral line shapes", "the line profile variations", "s", "oxygen abundance", "H α line", "Elemental abundances", "elemental abundances", "unusual composition", "Standard 1D local thermodynamic equilibrium analysis", "HD", "SUB", "Fe/H", "22.0 ±", "lt;italic&gt;T" ]
8.048456
9.376155
-1
12409132
[ "Mason, James Paul", "Woods, Thomas N.", "Webb, David F.", "Thompson, Barbara J.", "Colaninno, Robin C.", "Vourlidas, Angelos" ]
2016ApJ...830...20M
[ "Relationship of EUV Irradiance Coronal Dimming Slope and Depth to Coronal Mass Ejection Speed and Mass" ]
55
[ "Laboratory for Atmospheric and Space Physics, University of Colorado at Boulder, 3665 Discovery Drive, Boulder, CO 80303, USA", "Laboratory for Atmospheric and Space Physics, University of Colorado at Boulder, 3665 Discovery Drive, Boulder, CO 80303, USA", "Institute for Scientific Research, Boston College, Newton, MA 02458, USA", "NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Space Science Division, Naval Research Laboratory, Washington, DC 20009-1231, USA", "The Johns Hopkins University Applied Physics Laboratory, Laurel, MD 20732, USA" ]
[ "2017ApJ...836...17A", "2017JGRA..122.7979B", "2017JSWSC...7A...7D", "2017MNRAS.472..876O", "2018ApJ...855..137D", "2018ApJ...863..169D", "2018CQGra..35r5008P", "2018JHEAp..17....1P", "2018PhDT.......144D", "2019ApJ...874..123D", "2019ApJ...877..105M", "2019ApJ...884L..13A", "2019ApJS..244...13M", "2019BAAS...51c.300Y", "2019EPJC...79..288P", "2019MNRAS.486.3488A", "2020A&A...633A.142Z", "2020ApJ...889..104S", "2020ApJ...896...17C", "2020ApJ...905...17P", "2020FrASS...7...43V", "2020IAUS..354..426J", "2020IEEEA...8t0237A", "2020SpWea..1802382K", "2021ARA&A..59..445H", "2021GMS...258..179T", "2021LRSP...18....3V", "2021NatAs...5..697V", "2021PEPS....8...47W", "2021PEPS....8...56Z", "2021SSRv..217...82N", "2022ApJ...928..147A", "2022ApJ...928..154J", "2022ApJ...931...76X", "2022ApJ...933...92C", "2022ApJ...936..170L", "2022JATIS...8a4006F", "2022SerAJ.205....1L", "2022SoPh..297...59K", "2022Univ....8..565S", "2022arXiv221105506N", "2023A&A...678A.166C", "2023AJ....165...63F", "2023ApJ...949...61S", "2023ApJ...956...24K", "2023BAAS...55c.254L", "2023ChJSS..43..406B", "2023ScSnT..53.2021T", "2023arXiv231113942J", "2024A&A...683A..15J", "2024ApJ...961...23N", "2024ApJ...966...24Y", "2024Ge&Ae..64....1V", "2024arXiv240513671X", "2024arXiv240608194X" ]
[ "astronomy" ]
6
[ "methods: data analysis", "Sun: activity", "Sun: corona", "Sun: coronal mass ejections: CMEs", "Sun: flares", "Sun: UV radiation", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.3847/0004-637X/830/1/20", "10.48550/arXiv.1607.05284" ]
1607
1607.05284_arXiv.txt
\label{sec:intro} Large regions of temporary dimming or darkening of preexisting solar coronal emission often accompany coronal mass ejections (CMEs) and may trace field lines opened during the CME. The plasma of the solar corona responds in a number of ways to an eruptive event. \citet{Mason2014} provide details about the physics behind coronal dimming and the observational effects to be considered during analysis. Therein, the case is made for two hypotheses: that the slope of deconvolved, extreme ultraviolet (EUV) dimming irradiance light curves should be directly proportional to CME speed, and similarly, that dimming depth should scale with CME mass. Dimming regions can be extensive, representing at least part of the ``base" of a CME and the mass and magnetic flux transported outward by it. In this paper, we use the methods of \citet{Mason2014} to isolate EUV irradiance dimming as observed by Solar Dynamics Observatory (SDO; \citealt{Pesnell2012}) Extreme Ultraviolet (EUV) Variability Experiment (EVE; \citealt{Woods2012}) due to mass loss and characterize its time series in terms of slope and magnitude. \citet{Mason2014} focused on a single event and thus a correlation between dimming and CME parameters could not be established. Here, we analyze 37 events, 17 of which are used to establish a relationship between dimming slope and depth to CME speed and mass, respectively. We also outline the physical derivation for these relationships and assess their consistency with the data. Extensive surveys of EUV images containing coronal dimming events and their relation to CMEs have been performed by \citet{Reinard2008, Reinard2009}. For their sample of ~100 dimming events, \citet{Reinard2008} found mean lifetimes of 8 hours, with most disappearing within a day. \citet{Reinard2009} studied CMEs with and without associated dimmings, finding that those with dimmings tended to be faster and more energetic. \citet{Bewsher2008} found a 55\% association rate of dimming events with CMEs, and conversely that 84\% of CME events exhibited dimming. The timescale for dimming development is typically several minutes to an hour. This is much faster than the radiative cooling time, which implies that the cause of the decreased emission is more dependent on density decrease than temperature change \citep{Hudson1996}. Studies have demonstrated that dimming regions can be a good indicator of the apparent base of the white light CME \citep{Thompson2000, Harrison2003, Zhukov2004}. Thus, dimmings are usually interpreted as mass depletions due to the loss or rapid expansion of the overlying corona \citep{Hudson1998, Harrison2000, Zhukov2004}. Many landmark studies have established that dimmings can contribute a large fraction of the mass to a CME \citep{Harrison2000, Harrison2003, Zhukov2004, Aschwanden2009}. An Earth-directed CME's potential geoeffectiveness is typically characterized by four values: its velocity, mass, and the magnitude and duration of the southward component of the magnetic field (B$_z$) impacting Earth. Typical CME forecasts provide a predicted Earth arrival time only, which chiefly depends on velocity. The current standard process for estimating velocity relies on sequential coronagraph images from the Solar and Heliospheric Observatory's (SOHO; \citealt{Domingo1995}) Large Angle Solar Coronagraph C2 and C3 (LASCO; \citealt{Brueckner1995}) and the Solar Terrestrial Relations Observatory's (STEREO; \citealt{Kaiser2007}) COR1 and COR2 coronagraphs \citep{Howard2008}. Analysis of coronagraph images to determine CME velocities and masses results in relatively large uncertainties of 30-50\% \citep{Vourlidas2000, Vourlidas2010a, Vourlidas2011a}. The velocity and mass measurements with the most uncertainty are for Earth-directed CMEs that are seen as halos in coronagraphs at or near Earth. For these CMEs, speed determination is significantly affected by projection on the plane-of-sky, and a large percentage of the mass can be hidden behind the instrument's occulter. Without observations of these CMEs from another viewpoint, such as STEREO, it is difficult to make an accurate measurement of the CME velocity and mass from the coronagraph observations. However, EUV dimmings associated with these CMEs are very well observed by instruments in Earth orbit. Standard plane-of-sky velocity estimates are made and cataloged by the Coordinated Data Analysis Workshops (CDAW) CME catalog \citep{Gopalswamy2009}, which use routinely produced SOHO/LASCO coronagraph images. The different views from SOHO/LASCO and STEREO/COR images can be used to better constrain the velocity, direction, and mass of CMEs (e.g., \citep{Colaninno2009}). Coronal dimming can also be studied with spatially-integrated (full-disk) irradiance measurements as demonstrated by \citet{Mason2014}. They showed that a solar flare's impulsive and gradual phase peaks can initially dominate the irradiance for dimming-sensitive lines (e.g., Fe IX 171 \AA, Fe XII 195 \AA). They developed a technique for removing these flare peaks that only requires an independent, simultaneous irradiance measurement from a dimming-insensitive, flare-sensitive line (e.g. Fe XV 284 \AA), and they demonstrated how to apply this ``correction method" to the solar eruptive event on 2010 August 7. They also reported the kinetic energy parameters (mass and speed) of the associated coronal mass ejection (CME) that could be related to the dimming results of depth and slope. This follow-on paper expands that prior work to several more events in order to study the relationship of the CME and coronal dimming parameters. Here, we analyze 37 coronal dimming events during two separate two-week periods during 2011 and search for the relationship between dimming and CME speed and mass, plus the event from the \citet{Mason2014} paper, for a total initial sample of 38 events. Of the events studied, 17 could be parameterized in both dimming with SDO/EVE data and in CME velocity from SOHO/LASCO and STEREO/COR observations. 14 of the events yielded valid results in terms of dimming with SDO/EVE data and CME mass derived from the coronagraph observations. Section 2 describes the method for selecting this sample of events and explains why some events initially identified in SDO's Atmospheric Imaging Array (AIA; \citealt{Lemen2012}) could not be analyzed with SDO/EVE and/or STEREO/COR. Section 3 provides statistics on the flare-peak correction method detailed in \citet{Mason2014}, specifically which combinations of dimming and non-dimming lines provided the best correction for each of the events. Section 4 describes the fitting method applied to the corrected SDO/EVE light curves, including a discussion of uncertainties. Finally, Section 5 shows the correlations between the various combinations of coronal dimming and CME parameters, and conclusions about dimming and CME relationships are presented in Section 6.
Positive correlations with a high degree of significance have been found between coronal dimming and CME parameters. Our physically-motivated hypothesis that the CME mass goes as $\rm \sqrt{depth}$ had the second highest correlation and the scatterplots looked good when coronagraph-based masses below $10^{15}$ g were ignored. The second hypothesis, that CME speed should go as dimming slope/depth, was proven incorrect (barring a very unlikely sampling of the statistical space). However, the direct relationship between CME speed and dimming slope had a strong Pearson correlation coefficient and strong significance, though the scatterplot showed that there is a need for more data points. Future work will include hundreds to thousands of events, which should alleviate any concerns about statistical significance. Nevertheless, tentative equations relating CME mass and speed to EUV irradiance dimming depth and slope have been established in Equations \ref{eq:massdepth} and \ref{eq:speedslope}. Additionally, we found that the Fe IX 171 \AA\ dimming corrected for the flare contributions using the Fe XV 284 \AA\ line provides the most accurate dimming results for the SDO/EVE data. We note that the uncertainties for coronagraph and dimming parameters are complimentary: there are smaller uncertainties for CME speed than dimming slope, and there are smaller uncertainties for dimming depth than CME mass.
16
7
1607.05284
Extreme ultraviolet (EUV) coronal dimmings are often observed in response to solar eruptive events. These phenomena can be generated via several different physical processes. For space weather, the most important of these is the temporary void left behind by a coronal mass ejection (CME). Massive, fast CMEs tend to leave behind a darker void that also usually corresponds to minimum irradiance for the cooler coronal emissions. If the dimming is associated with a solar flare, as is often the case, the flare component of the irradiance light curve in the cooler coronal emission can be isolated and removed using simultaneous measurements of warmer coronal lines. We apply this technique to 37 dimming events identified during two separate two-week periods in 2011 plus an event on 2010 August 7, analyzed in a previous paper to parameterize dimming in terms of depth and slope. We provide statistics on which combination of wavelengths worked best for the flare-removal method, describe the fitting methods applied to the dimming light curves, and compare the dimming parameters with corresponding CME parameters of mass and speed. The best linear relationships found are v CME km s ≈ 2.36 × 10 6 km % × s dim % s m CME [ g ] ≈ 2.59 × 10 15 g % × d dim [ % ] . These relationships could be used for space weather operations of estimating CME mass and speed using near-real-time irradiance dimming measurements.
false
[ "CME mass", "corresponding CME parameters", "warmer coronal lines", "CME", "solar eruptive events", "simultaneous measurements", "minimum irradiance", "mass", "speed", "slope", "near-real-time irradiance dimming measurements", "a coronal mass ejection", "the cooler coronal emission", "the cooler coronal emissions", "37 dimming events", "the dimming light curves", "the dimming parameters", "space weather operations", "several different physical processes", "depth" ]
13.282783
16.119833
2
12510163
[ "Younsi, Ziri", "Zhidenko, Alexander", "Rezzolla, Luciano", "Konoplya, Roman", "Mizuno, Yosuke" ]
2016PhRvD..94h4025Y
[ "New method for shadow calculations: Application to parametrized axisymmetric black holes" ]
260
[ "Institute for Theoretical Physics, Goethe-University, Max-von-Laue-Str. 1, 60438 Frankfurt, Germany", "Centro de Matemática, Computação e Cognição, Universidade Federal do ABC (UFABC), Rua Abolição, CEP: 09210-180 Santo André, SP, Brazil; Institute for Theoretical Physics, Goethe-University, Max-von-Laue-Str. 1, 60438 Frankfurt, Germany", "Institute for Theoretical Physics, Goethe-University, Max-von-Laue-Str. 1, 60438 Frankfurt, Germany; Frankfurt Institute for Advanced Studies, Ruth-Moufang-Str. 1, D-60438 Frankfurt am Main, Germany", "Institute for Theoretical Physics, Goethe-University, Max-von-Laue-Str. 1, 60438 Frankfurt, Germany", "Institute for Theoretical Physics, Goethe-University, Max-von-Laue-Str. 1, 60438 Frankfurt, Germany" ]
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"2024arXiv240508062G", "2024arXiv240600579L" ]
[ "astronomy", "physics" ]
19
[ "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.1103/PhysRevD.94.084025", "10.48550/arXiv.1607.05767" ]
1607
1607.05767_arXiv.txt
It is now widely believed that at the centre of every galaxy resides a supermassive black hole. Observational evidence, particularly for our own Galactic black hole candidate, Sagittarius A* (Sgr A*), is compelling \cite{Eckart:1996,Gillessen:2009} and supports the notion of an object of enormous density, most likely a supermassive black hole, residing in the innermost central region. However, direct observation of an astrophysical black hole remains illusive, and this is because of the existence of the event horizon, that is, a surface limiting a region of spacetime beyond which neither matter nor radiation can escape the gravity of the black hole. Outside this surface, but still in close proximity to the event horizon, lies the photon-capture region, where photons follow unstable orbits. Hence, when observing a black hole directly, we expect to see a ``silhouette'' of this photon region. Therefore, black holes are expected to be observed as a ``shadow'' on the background sky \cite{Grenzebach:2014, Grenzebach:2016,Cunningham:1973}. It is anticipated that submillimeter very long baseline interferometry (VLBI) observations of Sgr A* with the Event Horizon Telescope (EHT) \cite{Huang:2007,Doeleman:2008,Goddi:2016} will soon yield the first radio images of the shadow of the candidate black hole therein. The Black Hole Camera project, in addition to other scientific activities, participates actively to the investigation of the physics and astrophysics of the black-hole candidate associated to Sgr A*. Particular attention is dedicated to theoretical calculations of the shadows, whose size and shape are sensitive to certain system parameters, in particular the black hole mass and spin, as well as the orientation of the spin axis of the black hole with respect to Earth (see, \eg Ref. \cite{Falcke:2000}). Observations of this shadow would not only provide very compelling evidence for the existence of an event horizon, but also enable estimates to be placed on these system parameters. Whilst astrophysical black holes are expected to be described by the Kerr solution, there exist numerous black hole solutions in other theories of gravity (see, for example, \cite{Stein:2014xba, Ayzenberg:2014aka} and references therein). One cannot yet exclude the possibility of many of the black hole metrics available in the literature and as such they are all, in a sense, potential candidates. Rather than investigate all possible theories of gravity and their corresponding black hole solutions one at a time, it is expedient to instead consider a model-independent framework within which any particular solution to any theory of gravity may be parameterised through a finite number of modifiable parameters. These parameters can then be chosen to measure deviations from the Kerr metric and may be estimated from astrophysical observations \cite{Vigeland:2011ji}. There is a simple reason why this avenue is a viable one, and although it is quite obvious, it may be useful to recall it here. The problem of defining the properties of the shadow does not require the choice of a theory of gravity, but only of a well-behaved expression for the metric tensor. This is because all that is ultimately needed to compute a shadow is the solution of the geodesic equations. The latter obviously do not require any assumption on the theory of gravity, but only a well defined and regular definition of the metric tensor. In Ref. \cite{Rezzolla:2014}, such a parametric framework was introduced to describe the spacetime of spherically symmetric and slowly rotating black holes in generic metric theories of gravity. The parameterisation in \cite{Rezzolla:2014} is based on a continued fraction expansion in terms of a compactified radial coordinate. Building upon the framework of \cite{Rezzolla:2014} to also include axisymmetric spacetimes, Ref. \citep{Konoplya:2016jvv} presented a parametric description of axisymmetric black holes in generic metric theories of gravity. This new parameterisation is based on a double expansion in both the radial and polar directions of a general stationary and axisymmetric metric, and is practically independent of any specific metric theory of gravity. Although it was shown to accurately reproduce, with only a small number of parameters, several different spacetime geometries, the question of how many expansion orders in each direction are required to accurately describe physical processes within this parametric framework was not addressed in Ref. \citep{Konoplya:2016jvv}. However, it is important to establish whether such a framework can reproduce, to high precision, the strong field behaviour of geodesics in the vicinity of the event horizon of different black hole spacetimes. Calculating the black hole shadow through direct numerical integration of the geodesic equations in the parameterised form of a reference black hole metric (and repeating the calculation at successive expansion orders), and subsequently comparing this with the shadow obtained from the analytic form of the ``un-parameterised" metric, provides a practical and stringent test of this framework. In addition, ray-tracing calculations of the shadows cast by different black hole solutions can provide insight into the practical performance of the application of this parameterisation in astrophysical calculations involving electromagnetic radiation. Such a framework has several important applications: \begin{enumerate} \item To enable black hole solutions in many metric theories of gravity to actually be written in algebraic form and therefore investigated using ray-tracing and radiative-transfer methods. \item To represent all black hole solutions in terms of just a few parameters, distinguishing between solutions on this basis. \item To constrain and potentially (physically) exclude black hole solutions from many theories of gravity with just a few key observational parameters necessary to reproduce the shadow curve (see \cite{Abdujabbarov:2015} for a general approach). \end{enumerate} In this first study we concern ourselves only with the shadow images obtained from black holes in different metric theories of gravity. Since the observed properties of radiation emanating from a black hole are subject to the spacetime through which the radiation propagates, it is prudent to first develop a method to ray-trace through a general parameterised metric and investigate the accuracy of this parameterisation. Hence, we here numerically calculate the shadow boundary curve and investigate, for several different spacetimes, the accuracy of the parameterisation at various orders with respect to the original un-parameterised form of the spacetime. Since the parameterisation exactly reproduces Kerr in the equatorial plane, and in order to adequately test the parameterisation, we consider near-extremal values of all spacetime-specific deformation parameters. Different measures of the accuracy of the expansion for each spacetime are presented and the excellent convergence properties of the parameterisation are demonstrated. This paper is organised as follows. In Sec. \ref{sec:rtf} we describe the ray-tracing formalism required to calculate geodesics within an arbitrary metric parameterisation, where the expressions for such calculations are derived explicitly. Section~ \ref{sec:pf} presents a short overview of the axisymmetric parameterisation framework employed throughout this paper. In Sec. \ref{sec:Kerr_Test} we apply this ray-tracing formalism to several different known black hole solutions. Each parameterised black hole solution is expanded to various orders and the resultant black hole shadows are calculated and compared with the shadow from the ``exact" metric. Finally, Sec. \ref{sec:c} is devoted to the conclusions.
\label{sec:c} We have introduced and subsequently employed a new method for performing general-relativistic ray-tracing calculations in order to calculate the black hole shadow images from a new parameterisation of any axisymmetric black hole metric. This new parameterisation can, with a small number of terms, represent any general stationary and axisymmetric black hole in any metric theory of gravity. We investigated and verified the effectiveness of this parameterisation for successive orders in the expansion, demonstrating both its convergence and accuracy for three different spacetimes: \begin{enumerate} \item The Kerr-Sen metric, fixed at second order in the polar expansion and varied from first through to fourth order in the radial expansion (the Kerr black hole is exactly reproduced at second order in the polar direction). \item A regular form of the EDGB metric, itself obtained from an expansion in terms of the parameters $\chi$ and $\zeta$. This regular solution is approximate but converges, to any desired accuracy, to the original approximate EDGB solution that diverges at $r=2M$. The expansion was purely polar and varied from first through to fourth order (in $\cos^{2}\theta$) in the polar direction. \item The Johannsen-Psaltis metric, represented as a double expansion in both the polar and radial directions. The expansion was fixed at second order in the polar direction and varied from first to eighth order in the radial direction. \end{enumerate} For all the aforementioned metrics, we chose values of the spin parameter and metric deformation parameters to be as extremal as possible whilst still ensuring the existence of an event horizon (\ie avoiding the appearance of a naked singularity). We performed three tests for each expansion order of each metric, calculating: (i) the black hole shadow polar curve obtained at each order, (ii) the error relative to the exact metric along the shadow curve, and (iii) the error of the half-area of the black hole shadow with respect to that obtained from the exact metric. Test (i) provided a qualitative comparison of the performance of the expansion as the order was increased, while test (ii) provided a quantification of the performance of the parameterisation everywhere along the shadow boundary. This test, in particular, provided an understanding of how well the parameterisation represents the spacetime, for example, in the equatorial plane and at the poles. Finally, test (iii) verified the excellent convergence behaviour and accuracy of the parameterisation as the expansion order was increased. We demonstrated that by increasing the order of polar and radial expansions the spacetime under consideration can be represented to essentially any desired accuracy\footnote{\texttt{Mathematica} notebooks containing the relevant coefficients for the metrics considered in this study may be found in the supplementary material accompanying this paper.}. Accurate calculations of black hole shadows in parameterised metrics represent a stringent test of parameterised representations of metric theories of gravity. Photons which delineate the shadow boundary pass very close to the event horizon and are subject to the steepest gradients of the gravitational potentials. Hence, accurately reproducing the behaviour of the spacetime in these regions lends credence to the prospect of employing this parameterisation framework to investigate not only black hole solutions in other metric theories of gravity, but to also perform the detailed radiative transport calculations required to investigate physical processes in other theories of gravity. Such calculations will prove useful for the interpretation of upcoming sub-mm VLBI observations from Sgr A* and for testing the Kerr black hole hypothesis. \begin{table*} \setlength{\tabcolsep}{7pt} \vspace*{-1mm} \caption{Table of the half-shadow area, $A_{n,m}$ (in units of $M^{2}$), and its corresponding percentage error (with respect to the analytic Kerr-Sen metric), $\epsilon_{n,m}$, of each expansion order $\{n,m\}$ of the Kerr-Sen metric. Indices $n$ and $m$ correspond, respectively, to the order of the \textit{radial} and $\cos\theta$ expansions. The exact area obtained from the analytic metric is denoted by $A_{\mathrm{exact}}$. Numbers within square brackets denote multiplicative powers of 10. For the Kerr-Sen metric $\cos\theta$ is fixed at second order and the \textit{radial} expansion is varied up to fourth order.} \centering \vspace{1mm} \begin{tabular}{| c | c | c | l l l l | l l l l |} % \hline % $a$ & $b$ & ${A_{\mathrm{\mathbf{exact}}}}$ & ${A_{1,2}}$ & ${A_{2,2}}$ & ${A_{3,2}}$ & ${A_{4,2}}$ & ${\epsilon_{1,2}}$ & ${\epsilon_{2,2}}$ & ${\epsilon_{3,2}}$ & ${\epsilon_{4,2}}$ \\ \hline \multirow{2}*{$0.2$} & $0.5$ & $73.5640$ & $73.5641$ & $73.5642$ & $73.5640$ & $73.5640$ & $1.06[-4]$ & $3.22[-4]$ & $3.08[-5]$ & $4.85[-6]$ \\[0.0ex] \cline{2-11} & $1.0$ & $110.576$ & $110.601$ & $110.581$ & $110.575$ & $110.576$ & $2.32[-2]$ & $4.38[-3]$ & $4.29[-4]$ & $2.75[-6]$ \\ [0.0ex] \cline{1-11} \multirow{2}*{$0.5$} & $0.5$ & $72.6472$ & $72.6289$ & $72.6478$ & $72.6471$ & $72.6472$ & $2.52[-2]$ & $8.20[-4]$ & $1.29[-4]$ & $8.18[-6]$ \\[0.0ex] \cline{2-11} & $1.0$ & $109.237$ & $109.241$ & $109.240$ & $109.235$ & $109.237$ & $3.83[-3]$ & $2.91[-3]$ & $1.28[-3]$ & $6.11[-5]$ \\ [0.0ex] \cline{1-11} \multirow{2}*{$0.95$} & $0.5$ & $68.3328$ & $67.9290$ & $68.4379$ & $68.3317$ & $68.3328$ & $5.91[-1]$ & $1.54[-1]$ & $1.49[-3]$ & $1.10[-4]$ \\[0.0ex] \cline{2-11} & $1.0$ & $102.953$ & $102.512$ & $103.081$ & $102.940$ & $102.954$ & $4.28[-1]$ & $1.25[-1]$ & $1.25[-2]$ & $9.51[-4]$ \\ [0.0ex] \cline{1-11} \multirow{2}*{$0.998$} & $0.5$ & $66.7159$ & $65.8826$ & $67.1017$ & $66.7136$ & $66.7160$ & $1.25[+0]$ & $5.78[-1]$ & $3.41[-3]$ & $1.90[-4]$ \\[0.0ex] \cline{2-11} & $1.0$ & $100.620$ & $99.7024$ & $101.032$ & $100.588$ & $100.622$ & $9.12[-1]$ & $4.10[-1]$ & $3.13[-2]$ & $1.77[-3]$ \\[0.0ex] \hline \end{tabular} \label{table:Sen} \end{table*} \begin{table*} \setlength{\tabcolsep}{7pt} \vspace*{0mm} \caption{Table of coefficients for the non-divergent EDGB metric.} % \centering \vspace{1mm} \begin{tabular}{| c | c | c | c c c | c c c |} % \hline ${k}$ & ${c_{k}}$ & ${w_{k}}$ & ${f_{k,1}}$ & ${f_{k,2}}$ & ${f_{k,3}}$ & ${\beta_{k,1}}$ & ${\beta_{k,2}}$ & ${\beta_{k,3}}$ \\ \hline $0$ & $-{4463}/{875}$ & $-9$ & ${3019}/{875}$ & $-{3048}/{875}$ & ${26}/{3}$ & $2$ & $5$ & $-14$ \\ \cline{1-9} $1$ & $-{2074}/{175}$ & $-140$ & ${11201}/{1750}$ & $-{18551}/{5250}$ & ${22}/{5}$ & $11$ & ${139}/{15}$ & $-{128}/{5}$ \\ [0.0ex] \cline{1-9} $2$ & $-{266911}/{12250}$ & $-90$ & $-{1497089}/{36750}$ & ${838039}/{110250}$ & ${32}/{5}$ & ${2767}/{15}$ & $-{907}/{45}$ & $-48$ \\[0.0ex] \cline{1-9} $3$ & $-{12673}/{525}$ & $-144$ & ${30316}/{3675}$ & $-{253756}/{11025}$ & $-{80}/{3}$ & $-{208}/{5}$ & ${616}/{5}$ & 0 \\[0.0ex] \cline{1-9} $4$ & ${12371}/{245}$ & $400$ & $-{26233}/{245}$ & ${1917}/{245}$ & 0 & ${1658}/{5}$ & ${2102}/{15}$ & $0$ \\[0.0ex] \cline{1-9} $5$ & ${3254}/{35}$ & --- & ${9214}/{21}$ & $-{20422}/{315}$ & $0$ & $-{26288}/{15}$ & ${28688}/{45}$ & $0$ \\[0.0ex] \cline{1-2}\cline{3-9} $6$ & ${2536}/{15}$ & --- & $-{6136}/{15}$ & ${7336}/{45}$ & $0$ & $2160$ & $-720$ & $0$ \\[0.0ex] \cline{1-2}\cline{3-9} $7$ & $-240$ & --- & $240$ & $-80$ & $0$ & --- & --- & --- \\[0.0ex] \hline \end{tabular} \label{table:coefficients} \end{table*} \begin{table*} \setlength{\tabcolsep}{7pt} \vspace*{-2mm} \caption{Table of half-shadow areas and their corresponding percentage errors for the first four successive expansion orders of the EDGB metric. The expansion considered in this case is in terms of $\cos^{2}\theta$ only.} \centering \vspace{1mm} \begin{tabular}{| c | c | c | l l l l | l l l l |} \hline $a$ & $\zeta$ & ${A_{\mathrm{exact}}}$ & ${A_{0,2}}$ & ${A_{0,4}}$ & ${A_{0,6}}$ & ${A_{0,8}}$ & ${\epsilon_{0,2}}$ & ${\epsilon_{0,4}}$ & ${\epsilon_{0,6}}$ & ${\epsilon_{0,8}}$ \\ \hline \multirow{2}*{$0.5$} & $0.1$ & $39.9885$ & $40.3445$ & $40.2700$ & $40.2559$ & $40.2088$ & $8.90[-1]$ & $7.04[-1]$ & $6.69[-1]$ & $5.51[-1]$ \\ \cline{2-11} & $0.15$ & $38.9264$ & $39.4915$ & $39.3773$ & $39.3536$ & $39.2825$ & $1.45[+0]$ & $1.16[+0]$ & $1.10[+0]$ & $9.15[-1]$ \\ \hline \end{tabular} \label{table:EDGB} \end{table*} \begin{table*} \setlength{\tabcolsep}{5pt} \vspace*{-2mm} \centering \caption{Table of half-shadow areas for successive expansion orders of the Johannsen-Psaltis metric. The expansion is fixed at fourth order for $\cos\theta$ whilst the \textit{radial} expansion is considered up to eighth order. Note that for the eighth order radial expansion the value of $A_{8,2}$ (\ie second order in $\cos\theta$) is included to compare with $A_{8,4}$. } \centering \vspace{1mm} \begin{tabular}{| c | c | c | l l l l l l l l l |} \hline \vspace*{-1mm} $a$ & $\varepsilon_{3}$ & ${A_{\mathrm{exact}}}$ & ${A_{1,4}}$ & ${A_{2,4}}$ & ${A_{3,4}}$ & ${A_{4,4}}$ & ${A_{5,4}}$ & ${A_{6,4}}$ & ${A_{7,4}}$ & ${A_{8,2}}$ & ${A_{8,4}}$ \\ & & & & & & & & & & & \\ [-1ex] \cline{4-12} \hline \\ [-3.15ex] \multirow{3}*{$0.9$} & \multirow{1}*{$0.24$} & \multirow{1}*{$38.8979$} & $39.5424$ & $38.5721$ & $38.9766$ & $38.8723$ & $38.9158$ & $38.9175$ & $38.8838$ & $38.9171$ & $38.8888$ \\[0.0ex] \cline{4-12} \\ [-3.15ex] \cline{2-12} \\ [-3.15ex] & \multirow{1}*{$-0.5$} & \multirow{1}*{$41.2918$} & $40.4947$ & $41.0235$ & $41.2720$ & $41.1276$ & $41.3215$ & $41.3076$ & $41.2778$ & $41.2747$ & $41.2880$ \\[0.0ex] \cline{2-12} \\ [-3.15ex] & \multirow{1}*{$-1.0$} & \multirow{1}*{$42.5892$} & $41.5207$ & $41.1449$ & $42.5444$ & $42.4225$ & $42.6546$ & $42.5539$ & $42.5356$ & $42.5678$ & $42.5806$ \\[0.0ex] \hline \end{tabular} \label{table:JP_1} \end{table*} \begin{table*} \setlength{\tabcolsep}{5pt} \vspace*{-2mm} \centering \caption{Table of percentage errors corresponding to the half-shadow areas for successive expansion orders of the Johannsen-Psaltis metric as reported in Table \ref{table:JP_1}. As expected, $\epsilon_{8,4} < \epsilon_{8,2}$ for all values of $\varepsilon_{3}$. } \centering \vspace{1mm} \begin{tabular}{| c | c | c l l l l l l l l |} \hline \vspace*{-1mm} $a$ & $\varepsilon_{3}$ & ${\epsilon_{1,4}}$ & ${\epsilon_{2,4}}$ & ${\epsilon_{3,4}}$ & ${\epsilon_{4,4}}$ & ${\epsilon_{5,4}}$ & ${\epsilon_{6,4}}$ & ${\epsilon_{7,4}}$ & ${\epsilon_{8,2}}$ & ${\epsilon_{8,4}}$ \\ [1ex] \hline \\ [-3.15ex] \multirow{3}*{$0.9$} & \multirow{1}*{$0.24$} & $1.66[+0]$ & $8.38[-1]$ & $2.02[-1]$ & $6.59[-2]$ & $4.58[-2]$ & $5.02[-2]$ & $3.63[-2]$ & $4.91[-2]$ & $2.34[-2]$ \\ [0.0ex] & \multirow{1}*{$-0.5$} & $1.93[+0]$ & $6.50[-1]$ & $4.78[-2]$ & $3.98[-1]$ & $7.19[-2]$ & $3.82[-2]$ & $3.39[-2]$ & $4.13[-2]$ & $9.20[-3]$ \\ [0.0ex] & \multirow{1}*{$-1.0$} & $2.51[+0]$ & $3.39[+0]$ & $1.05[-1]$ & $3.91[-1]$ & $1.54[-1]$ & $8.29[-2]$ & $1.26[-1]$ & $5.02[-2]$ & $2.00[-2]$ \\ [0.0ex] \hline \end{tabular} \label{table:JP_2} \end{table*}
16
7
1607.05767
Collaborative international efforts under the name of the Event Horizon Telescope project, using sub-mm very long baseline interferometry, are soon expected to provide the first images of the shadow cast by the candidate supermassive black hole in our Galactic center, Sagittarius A*. Observations of this shadow would provide direct evidence of the existence of astrophysical black holes. Although it is expected that astrophysical black holes are described by the axisymmetric Kerr solution, there also exist many other black hole solutions, both in general relativity and in other theories of gravity, which cannot presently be ruled out. To this end, we present calculations of black hole shadow images from various metric theories of gravity as described by our recent work on a general parametrization of axisymmetric black holes [R. Konoplya, L. Rezzolla, and A. Zhidenko, Phys. Rev. D 93, 064015 (2016).]. An algorithm to perform general ray-tracing calculations for any metric theory of gravity is first outlined and then employed to demonstrate that even for extremal metric deformation parameters of various black hole spacetimes, this parametrization is both robust and rapidly convergent to the correct solution.
false
[ "many other black hole solutions", "axisymmetric black holes", "black hole shadow images", "astrophysical black holes", "various black hole spacetimes", "various metric theories", "other theories", "extremal metric deformation parameters", "general relativity", "the candidate supermassive black hole", "gravity", "Sagittarius A", "the axisymmetric Kerr solution", "sub-mm very long baseline interferometry", "the correct solution", "Phys", "general ray-tracing calculations", "direct evidence", "any metric theory", "A. Zhidenko" ]
8.869122
2.959978
61
12464596
[ "Kouvaris, Chris", "Langæble, Kasper", "Grønlund Nielsen, Niklas" ]
2016JCAP...10..012K
[ "The spectrum of darkonium in the Sun" ]
27
[ "CP3-Origins, Centre for Cosmology and Particle Physics Phenomenology, University of Southern Denmark, Campusvej 55, Odense M, 5230 Denmark", "CP3-Origins, Centre for Cosmology and Particle Physics Phenomenology, University of Southern Denmark, Campusvej 55, Odense M, 5230 Denmark", "CP3-Origins, Centre for Cosmology and Particle Physics Phenomenology, University of Southern Denmark, Campusvej 55, Odense M, 5230 Denmark" ]
[ "2017JCAP...02..005A", "2017JCAP...05..036C", "2017JHEP...02..091L", "2017JHEP...04..077P", "2017JHEP...11..108B", "2017PhRvD..95g5015S", "2017PhRvD..95l3016L", "2017PhRvD..96e5027L", "2018JHEP...07..096H", "2018JHEP...11..084B", "2018PhR...730....1T", "2018PhRvD..98l3012A", "2019JCAP...11..011N", "2019JHEP...01..070O", "2019JHEP...04..130H", "2019arXiv191106788B", "2020Galax...8...42S", "2020PhRvD.101e5044B", "2020ScPP....9...68B", "2020arXiv201213170D", "2021PhRvD.103j3005S", "2021PhRvD.104b3024B", "2022EPJC...82..142D", "2022JHEP...10..186L", "2022PhLB..83337323B", "2024JHEP...05..045C", "2024MNRAS.527.4483D" ]
[ "astronomy", "physics" ]
2
[ "High Energy Physics - Phenomenology", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1949PhRv...75.1696O", "1957qmot.book.....B", "1971rqt..book.....B", "1975PhR....22..215S", "1982ZhETF..83..892O", "1985ApJ...296..679P", "1985PhRvL..55..257S", "1986PhLB..166..196H", "1986PhLB..167..295F", "1986PhRvD..33.2079K", "1986PhRvD..34.2206G", "1987ApJ...321..560G", "1987ApJ...321..571G", "1987NuPhB.283..681G", "1993MNRAS.264..201K", "1993ZPhyC..60..175G", "1993ppc..book.....P", "1994ApJ...427L...1F", "1994Natur.370..629M", "1997ApJ...490..493N", "1998PhLB..432..134C", "1999ApJ...522...82K", "1999ApJ...524L..19M", "2003PhRvL..90v5002K", "2008PhLB..662...53P", "2009ApJ...705L.123S", "2009JCAP...07..004F", "2009MNRAS.397.1169D", "2010Ap&SS.328...13S", "2010JCAP...02..028S", "2010JHEP...06..029M", "2010JPhG...37j5009C", "2010PhRvD..81a6002S", "2010PhRvD..81j4019C", "2010PhRvD..82h3525F", "2010PhRvD..82k5012S", "2011ApJ...733...62L", "2011ApJ...738..102T", "2011MNRAS.415L..40B", "2011PhRvD..84c2007A", "2012ApJ...749...75S", "2012JCAP...07..043I", "2013MNRAS.430...81R", "2013MNRAS.430..105P", "2013PhLB..725..190A", "2013PhRvD..87k5007T", "2013PhRvD..87l3531P", "2013PhRvL.111d1302A", "2014JCAP...03..053B", "2014JCAP...12..033V", "2014PhRvD..89j3510L", "2014PhRvD..90c5022F", "2014PhRvE..89c2145K", "2014PhRvL.112i1303A", "2015JCAP...10..055D", "2015JHEP...07..045B", "2015JPhCS.587a2028G", "2015PhLB..747..331A", "2015PhRvD..92h3509L", "2016PhRvA..93a2506A", "2016PhRvD..93a5014F", "2016PhRvD..93b3527S", "2016PhRvD..93k5020A", "2016PhRvD..93k5036F", "2016PhRvD..94j3003I", "2016PhRvL.116o1801A" ]
[ "10.1088/1475-7516/2016/10/012", "10.48550/arXiv.1607.00374" ]
1607
1607.00374_arXiv.txt
Dark matter (DM) is approximately five times more abundant than baryonic matter in the Universe. Although DM can be in the form of more conventional compact objects like primordial black holes~\cite{Carr:2009jm}, there is also the possibility that DM might be in the form of particles. No Standard Model (SM) particle can play the role of DM. Therefore, in the case DM is in the form of particles, it must be related to physics beyond the SM. The simplest example of such a realization is the Weakly Interacting Massive Particle (WIMP) paradigm. In that case, WIMPs are produced in the early Universe and occasionally annihilate to SM particles. If the annihilation cross section is appropriate, annihilations are sufficient to produce the DM abundance we observe today. Such a scenario can potentially be tested. Annihilation of WIMPs in the Sun and the center of the Galaxy to SM particles could create observable signals. Similarly, this scenario allows production of WIMPs in collider experiments as long as there is enough energy to produce them or create nuclear recoils when WIMPs scatter off nuclei in underground detectors. A minimal extension of the SM that could facilitate the above characteristics can be realized by adding a new $U(1)$ gauge symmetry which breaks spontaneously providing the dark photon with a mass~\cite{Kobzarev:1966qya,Okun:1982xi}. In this scenario, the DM particle is charged under the $U(1)$ symmetry and since it is possible to have a small kinetic mixing between the dark photon and the SM photon under very generic grounds \cite{Holdom:1985ag}, the dark and bright sectors are linked. As we mentioned above, a particular way of testing this scenario is by searching for signals of DM annihilation in the Sun. DM particles can be trapped in the Sun simply by interacting with nuclei or electrons as they cross it. Trapped DM particles may thermalize with the interior of the Sun and sink to the center where they can meet each other and annihilate to SM particles. In particular, annihilation to neutrinos (which can easily escape from the Sun) may potentially lead to detectable signals in Earth-based detectors. The capture and annihilation process in the Earth and the Sun as well as the produced neutrino spectrum from DM annihilation has been studied extensively in the past~\cite{Press:1985ug,Gould:1987ir,Gould:1987ju,Silk:1985ax,Freese:1985qw,Krauss:1985aaa,Griest:1986yu}. Within the context of dark photons, similar studies have also been made regarding indirect signals from the Sun and the Earth or direct detection~\cite{Pospelov:2007mp,Schuster:2009au,Schuster:2009fc,Meade:2009mu,An:2013yfc,An:2013yua,Fradette:2014sza,An:2014twa,Feng:2015hja, Feng:2016ijc}. In this paper we investigate fermionic DM that self-interacts via a light dark photon mediator. In particular, we are interested in the region of the parameter space, where DM captured by the Sun has a substantial probability of recombining to positronium-like bound states of DM and anti-DM, which we from now on call darkonium. Darkonium states can subsequently decay to either two dark photons in the case of para-darkonium (where the spin of DM and anti-DM are opposite) or to three dark photons in the case of ortho-darkonium (where the spin of DM and anti-DM are aligned). We assume that dark photons are linked to SM via kinetic mixing with the ordinary photons. At the end of the day, the spectrum of the produced dark photons in the Sun has four components: i) monochromatic dark photons of energy $m_X$ (the mass of DM) which come from direct DM annihilation ii) monochromatic dark photons of energy $m_X-\Delta/2$ where $\Delta$ is the binding energy of the para-darkonium iii) dark photons that span energy up to $m_X-\Delta/2$ from the decay of ortho-darkonium states and iv) monochromatic photons of energy $\Delta$ produced from the recombination of DM to darkonium. We investigate under what conditions the aforementioned dark photons leave the Sun without decaying to SM particles. Those that exit the Sun intact, will decay sooner or later producing an electron-positron pair due to the kinetic mixing. Based on that, we find the precise positron spectrum that potentially could be observed by AMS-02 and we show how one can set constraints on these models. Interestingly, we find that there is parameter space where recombination dark photons could be detected earlier than the annihilation ones. We should also mention at this point that the possibility of forming darkonium can also have important consequences in the physics of the early Universe. For instance, if DM is thermally produced by a freeze-out of DM annihilation into mediators, the effect of darkonium bound states is to delay the freeze-out until later times~\cite{vonHarling:2014kha}. Additionally, the effect of the bound states in indirect Galactic searches and on the Galactic structure has been investigated in~\cite{An:2016gad,Feng:2009mn}. The indirect detection signal in models with light dark photons in the center of the Earth and the Sun have been recently investigated in~\cite{Feng:2015hja,Feng:2016ijc} albeit neglecting DM bound states. As we will show, in the region of the parameter space where darkonium can form, the spectrum of positrons is dominated by the decay pattern of the bound states, and direct annihilation contributes a subleading part in the full signal. The paper is structured as follows: In section \ref{Sec: DM in the Sun} we review the DM model and the necessary formalism to calculate the solar DM and darkonium populations. In section \ref{Sec: Parameter space} we demarcate the allowed parameter space in which darknonium can lead to a significant indirect detection signal. In section \ref{Sec: Kinematics of mediator decay} we review the kinematics and geometry of the positron signal from a decaying mediator emitted from the Sun. In section \ref{Sec: Results} we present the positron spectra and compare them with the flux observed by AMS-02. In section \ref{Sec: Conclusions} we make our concluding remarks. Throughout the paper we use natural units, i.e. $\hbar = c = k_\text{B}= 1$.
\label{Sec: Conclusions} In this paper we studied the positron spectrum produced by decaying dark photons that originate in the Sun. We consider dark photons that are produced inside the Sun via direct annihilation of $X\bar{X}$, decay of bound states $X\bar{X}$ and bound state formation. These photons decay to positron-electron pairs and those decays that take place outside the Sun can potentially create a positron flux on Earth, that could be detected e.g. by AMS-02. We demonstrated that the spectrum has distinct features as a function of the positron energy that could distinguish it from any other astrophysical background of positrons. These features are due to contributions from different types of dark photons i.e. coming from direct annihilation, decay of para- or ortho-darkonium and recombination. More importantly we find that there is parameter space where AMS-02 or equivalent experiments would be able to pick up the contribution from the dark recombination photons before they have a chance to observe the positrons that come from dark photons produced by the typical direct annihilation of DM inside the Sun. This result appears to be surprising since in principle recombination photons are fewer and less energetic than those from direct annihilation. However as we have argued, the positrons produced from the decay of recombination photons span a much narrower energy band than the rest, leading to higher numbers in lower energy bins. Possible discovery of a positron spectrum with the morphology of our Fig.~\ref{Fig:SplitSignal} not only can associate beyond any doubt the positron signal to its DM origin, but it could establish the type and the mass of the mediator, thus understanding how DM self-interacts.
16
7
1607.00374
Dark matter that gets captured in the Sun may form positronium-like bound states if it self-interacts via light dark photons. In this case, dark matter can either annihilate to dark photons or recombine in bound states which subsequently also decay to dark photons. The fraction of the dark photons that leave the Sun without decaying to Standard Model particles have a characteristic energy spectrum which is a mixture of the direct annihilation process, the decays of ortho- and para- bound states and the recombination process. The ultimate decay of these dark photons to positron-electron pairs (via kinetic mixing) outside the Sun creates a distinct signal that can either identify or set strict constraints on dark photon models.
false
[ "dark photons", "light dark photons", "dark photon models", "bound states", "Dark matter", "dark matter", "Sun", "Standard Model particles", "strict constraints", "the dark photons", "these dark photons", "kinetic mixing", "positronium-like bound states", "Standard Model", "ortho-", "the direct annihilation process", "the recombination process", "a characteristic energy spectrum", "positron-electron pairs", "a distinct signal" ]
8.141178
-1.685976
54
1758494
[ "Ridolfi, A.", "Freire, P. C. C.", "Torne, P.", "Heinke, C. O.", "van den Berg, M.", "Jordan, C.", "Kramer, M.", "Bassa, C. G.", "Sarkissian, J.", "D'Amico, N.", "Lorimer, D.", "Camilo, F.", "Manchester, R. N.", "Lyne, A." ]
2016MNRAS.462.2918R
[ "Long-term observations of the pulsars in 47 Tucanae - I. A study of four elusive binary systems" ]
56
[ "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany; Department of Physics, University of Alberta, CCIS 4-183, Edmonton, AB T6G 2E1, Canada", "Anton Pannekoek Institute for Astronomy, University of Amsterdam, Science Park 904, NL-1098 XH Amsterdam, the Netherlands; Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, UK", "Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-53121 Bonn, Germany; Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, UK", "ASTRON, the Netherlands Institute for Radio Astronomy, Postbus 2, NL-7990 AA Dwingeloo, the Netherlands", "CSIRO Astronomy and Space Science, Australia Telescope National Facility, Box 76, Epping, NSW 1710, Australia", "Osservatorio Astronomico di Cagliari, INAF, via della Scienza 5, I-09047 Selargius (CA), Italy; Dipartimento di Fisica, Università degli Studi di Cagliari, SP Monserrato-Sestu km 0,7, I-90042 Monserrato (CA), Italy", "Department of Physics and Astronomy, West Virginia University, PO Box 6315, Morgantown, WV 26506, USA", "Square Kilometre Array South Africa, Pinelands 7405, South Africa", "CSIRO Astronomy and Space Science, Australia Telescope National Facility, Box 76, Epping, NSW 1710, Australia", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Manchester M13 9PL, UK" ]
[ "2017ApJ...839...12A", "2017IJMPD..2630021M", "2017MNRAS.467.2199B", "2017MNRAS.471..857F", "2017MNRAS.472.3706B", "2017Natur.542..203K", "2017mbhe.confE..51K", "2018IAUS..337..251R", "2018IAUS..337..295A", "2018MNRAS.473..897N", "2018MNRAS.476.4794F", "2018MNRAS.481..627A", "2018MNRAS.481.4397C", "2018PhRvD..98d1301B", "2019ApJ...875....1M", "2019ApJ...875...25C", "2019ApJ...885L..37Z", "2019MNRAS.483.4560Z", "2019MNRAS.487..769A", "2020ApJ...888L..18D", "2020ApJ...892L...6P", "2020ApJ...901...57B", "2020MNRAS.491..113H", "2020MNRAS.493.2171D", "2020MNRAS.493.6033Z", "2020NatAs...4..704A", "2020PhRvD.101f3016B", "2020arXiv200200978Z", "2021A&A...649A.120M", "2021ApJ...912..124B", "2021MNRAS.500.1139H", "2021MNRAS.502.1596Z", "2021MNRAS.504.1407R", "2021NewA...8501549P", "2021arXiv210614535S", "2022A&A...661A..35S", "2022A&A...663A..75W", "2022A&A...664A..54G", "2022ApJ...931...84Y", "2022MNRAS.511.5964Z", "2022MNRAS.516.5309V", "2023ApJ...942...87K", "2023ApJ...942L..40I", "2023ApJ...944..225L", "2023ApJ...945...70Z", "2023ApJ...956..132L", "2023ApJS..269...56Z", "2023MNRAS.518.1642A", "2023MNRAS.519.5590C", "2023MNRAS.520.3847C", "2023MNRAS.521.2616D", "2023MNRAS.525L..76H", "2023MNRAS.526.2736Z", "2024A&A...682A..22D", "2024A&A...686A.166P", "2024MNRAS.530.1436V" ]
[ "astronomy" ]
20
[ "binaries: general", "pulsars: individual: PSR J0024-7204P", "pulsars: individual: PSR J0024-7204V", "pulsars: individual: PSR J0024-7204W", "pulsars: individual: PSR J0024-7201X", "globular clusters: individual: 47 Tucanae", "Astrophysics - High Energy Astrophysical Phenomena" ]
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[ "10.1093/mnras/stw1850", "10.48550/arXiv.1607.07248" ]
1607
1607.07248_arXiv.txt
\label{sec:intro} \newcommand{\dmunit}{pc\,cm$^{-3}$} \newcommand{\msun}{M$_{\odot}$} \newcommand{\rsun}{R$_{\odot}$} \newcommand{\us}{$\mu$s} \newcommand{\TO}{$T_0$} \newcommand{\ergs}{~erg\,s$^{-1}$} \newcommand{\chisquare}{\chi_{\textrm{red}}} Globular clusters (GCs) are known to be host of a wealth of radio pulsars: with 146 currently known\footnote{\url{http://www.naic.edu/~pfreire/GCpsr.html}}, the pulsars in GCs account for more than 5\% of the total known pulsar population\footnote{\url{http://www.atnf.csiro.au/research/pulsar/psrcat} (version 1.54, \citealt{mht+05})} and a large fraction of millisecond pulsars (MSPs, here defined as those pulsars with spin periods of 10 ms or less). \textcolor{black}{The population of pulsars in GCs also sharply contrasts with that of our Galaxy. The main reason is related to the much higher stellar densities that can be found in GCs, especially in their cores, which exceed those in the Galactic disk by several orders of magnitude. Such crammed environments provide the ideal conditions for the dynamical interaction of stars in two-body or even three-body encounters, the latter being the case of a star gravitationally interacting with a binary system. In this continuous exchange of gravitational energy, the more massive objects tend to sink towards the center of the cluster, a phenomenon called mass segregation \citep[see e.g. ][]{hge+05}. The process thus greatly fosters a concentration of neutron stars (NS) near the GC core and, consequently, the formation of exotic binary systems in which the NS can be spun up (or ``recycled'') by accreting matter and angular momentum from the companion star \citep{acr+82,rs82}. The observational evidence for this process is striking, as over 80\% of the GC pulsars known are MSPs and most of them reside close to the center of their host cluster. Nevertheless, there are substantial differences in the populations of pulsars in different GCs, particularly in the ratio between the isolated and the binary pulsars. As \citet{vf14} discussed, this ratio seems to be related to the encounter rate per binary of the GC: the higher the rate, the larger the fraction of isolated pulsars we are likely to see in the cluster.} \textcolor{black}{47 Tucanae (also known as NGC 104 and, hereafter, 47~Tuc) is a well known GC visible in the southern sky. The previous observations of 47 Tuc with the Parkes radio telescope have provided outstanding scientific results that include the discovery of 25 radio MSPs \citep{mlj+90,mlr+91,rlm+95,clf+00,phl+16}, all of which have spin periods smaller than 8 ms. This is a very exceptional population, since only $\sim$\,7\% of the non-cluster pulsars in the Galaxy have spin periods of less than 8 ms. As expected from a GC with a relatively low encounter rate per binary \citep{vf14}, a large fraction (60\%) of the pulsars in 47 Tuc are found in binary systems. This is similar to that for the Galactic non-cluster MSP population ($\sim$\,67\%), but very different from the general pulsar population, where the number in binaries is $\sim\,$7\%. The binary pulsars in 47 Tuc include five ``Black Widow" pulsars (BWPs), with orbital periods of $P_b \sim 0.06 - 0.4$~days and companion masses of $M_{\rm c}~\sim0.03-0.05$~\msun, and at least two eclipsing ``Redback'' pulsars (RBs) with $P_b \sim 0.1 - 0.2$~days and $M_{\rm c} \sim0.1-0.5$~\msun (see e.g. \citealt{f05,r13} for reviews on BWPs and RBs)}. The two categories share several common features, such as very short orbital periods (of the order of a few hours), highly recycled pulsars, very low mass companions that are undergoing mass loss, and, very often, regular eclipses in the pulsar radio signal. However, while in BWPs the companion stars have extremely low masses ($M_{\rm c} \lesssim 0.1$~\msun) and are being heavily ablated by the strong pulsar wind (which can make them almost completely stripped stars), in RBs the companions are mostly heavier ($M_{\rm c} \sim 0.1-0.5$~\msun), non-degenerate stars and the mass loss is usually driven by a Roche-lobe overflow. Of all the 47 Tuc pulsars, 17 have published timing solutions \citep{rlm+95,clf+00,fcl+01,fck+03,phl+16}, which allowed a detailed study of the dynamics of the cluster \citep{fcl+01}, the first detection of ionized gas in a globular cluster \citep{fkl+01}, the detection of all the pulsars at X-ray wavelengths \citep{gch+02,bgv05,bgh+06} and the optical detection of 5 of the 8 known MSP$-$White-Dwarf (WD) systems \citep[\textcolor{black}{namely} 47 Tuc Q, S, T, U and Y;][]{egh+01,cpf+15,rvh+15}. \textcolor{black}{Since the previous publication of radio timing data by \citet{fck+03}}, we have continued timing the millisecond pulsars in 47 Tuc. The main objectives of this long-term timing programme are diverse, and include: a) avoiding the loss of phase-connection \textcolor{black}{in the timing} of any of the pulsars; b) monitoring orbital phase variations for the BWP and RB binaries; \textcolor{black}{this is particularly important not only to understand their unpredictable orbital dynamics, but also in relation to the previous point, since losing track of the precise orbital phase can easily cause a loss of phase-connection in the timing;} c) improving the orbital measurements for the stable MSP$-$WD binaries, in particular the measurement of relativistic and kinematic effects; d) improving the determination of the proper motions of all pulsars, in order to study the internal dynamics of the cluster; e) determining orbits and timing solutions for previously known binary pulsars that still lack them; f) discovering new pulsars. This is part of a series of papers that will present the results of this \textcolor{black}{long-term programme}. In the work of \citet{phl+16} and in the upcoming papers, we address the search for new pulsars. In \textcolor{black}{the present} paper, we have used 16 years of archival data taken at the Parkes radio telescope to further characterize four among the faintest binary systems of 47 Tuc, namely pulsars P, V, W and X, for which an accurate orbital solution (and, hence, also a timing solution) was not known. Accurate orbital and timing solutions for pulsars R and Y will instead be presented in Freire et al. (in preparation). In Section 2 we describe our dataset. In Sections 3 we explain the two search methods that we have used to re-detect the four pulsars and we give refined orbital parameters for all of them. In Section 4 we describe our timing analysis, whose results are presented \textcolor{black}{and discussed} in Section 5. In Section 6 we summarize our findings. \vskip 0 cm
\label{sec:conclusions} We have presented a study of four binary millisecond pulsars belonging to the globular cluster 47 Tucanae that, despite being known, had never been investigated in detail because of the sparsity of their detections. We have used different search techniques on 16 years of archival data, which allowed us to detect each pulsar at many more epochs. Thanks to these, we were able to time and thus better characterize the pulsars. For two pulsars, namely 47 Tuc P and V, the detection rate was still too low to allow us to obtain a phase-connected timing solution, but we were able to obtain very precise measurements of their spin periods, DM's and orbital parameters. 47 Tuc V is a Redback-like binary system, eclipsed for 50\% of its orbit. Although its characteristics and behaviour are quite similar to those of the so-called ``transitional" MSPs, we showed that it seems unlikely that 47 Tuc V actually belongs to the latter class. For 47 Tuc W and X, we successfully obtained phase-connected timing solutions that allowed a full characterization of both systems. The former is a well-known Redback that has already been extensively studied at other wavelengths. Our radio timing complemented those studies and revealed a strong orbital variability, as it is often the case for Redbacks. 47 Tuc X is probably the most peculiar system and stands out for its long orbital period, its extreme circularity, as well as its distance from the cluster center, compared to all the other pulsars in 47 Tuc. We have studied this system in the radio, X-ray and optical wavelengths and discussed its possible formation. \textcolor{black}{The results presented in this work allowed a more complete characterization of the pulsar population of 47 Tuc. The properties of this population largely conform to the expectations of \citet{vf14} for a cluster with a low rate of stellar interactions per binary. In such a cluster, exchange encounters form NS$-$low-mass MS star binaries; these are the only type of MS star remaining in any globular cluster. The slow evolution of such stars and resulting long (and, in the case of 47 Tuc, undisturbed) accretion episode then lead invariably to fast-spinning pulsars (as observed in this and other lower density clusters) in systems similar to the binary MSPs in the Galaxy: MSP$-$He WD systems, Redbacks, Black-Widows and (perhaps evolving from the latter) isolated MSPs. Indeed, none of the binaries found so far in 47 Tuc is clearly the result of an exchange interaction. The only possible exception to this is 47~Tuc~X, an unusual system that could either have resulted from a primordial binary, or could have formed in an exchange encounter from a population of neutron stars that had not yet been segregated into the core of the cluster.} In the last few years, new and more sensitive observations of 47 Tuc have been carried out with the Parkes radio telescope. Even though it is unlikely that these new data will allow us to obtain a phase-connected solution for 47 Tuc P in the short term, the chances that this could be the case for pulsar V look higher. If so, this could help confirm or reject the hypothesis of a non-transitional nature for the object. For 47 Tuc X, the extended data span could improve the measurement of the orbital decay and the proper motion. A possible major leap in the study of these four pulsars might soon be possible with the MeerKAT\footnote{\url{http://public.ska.ac.za/meerkat}} telescope \citep[][]{bj12}, which, thanks to its position in the southern hemisphere, will allow a study of 47 Tuc at radio wavelengths with an unprecedented sensitivity.
16
7
1607.07248
For the past couple of decades, the Parkes radio telescope has been regularly observing the millisecond pulsars in 47 Tucanae (47 Tuc). This long-term timing programme was designed to address a wide range of scientific issues related to these pulsars and the globular cluster where they are located. In this paper, the first of a series, we address one of these objectives: the characterization of four previously known binary pulsars for which no precise orbital parameters were known, namely 47 Tuc P, V, W and X (pulsars 47 Tuc R and Y are discussed elsewhere). We determined the previously unknown orbital parameters of 47 Tuc V and X and greatly improved those of 47 Tuc P and W. For pulsars W and X we obtained, for the first time, full coherent timing solutions across the whole data span, which allowed a much more detailed characterization of these systems. 47 Tuc W, a well-known tight eclipsing binary pulsar, exhibits a large orbital period variability, as expected for a system of its class. 47 Tuc X turns out to be in a wide, extremely circular, 10.9-d long binary orbit and its position is ∼3.8 arcmin away from the cluster centre, more than three times the distance of any other pulsar in 47 Tuc. These characteristics make 47 Tuc X a very different object with respect to the other pulsars of the cluster.
false
[ "pulsars", "X", "47 Tuc X", "any other pulsar", "the other pulsars", "W", "47 Tuc W", "47 Tuc V", "47 Tuc P", "the millisecond pulsars", "scientific issues", "47 Tuc R", "V", "four previously known binary pulsars", "a large orbital period variability", "(pulsars", "these pulsars", "Y", "no precise orbital parameters", "the cluster centre" ]
6.283942
4.731439
-1
12433257
[ "Aparicio, Luis", "Cicoli, Michele", "Dutta, Bhaskar", "Muia, Francesco", "Quevedo, Fernando" ]
2016arXiv160700004A
[ "Light Higgsino Dark Matter from Non-thermal Cosmology" ]
13
[ "-", "-", "-", "-", "-" ]
[ "2016JHEP...11..182C", "2017JHEP...08..072K", "2017JHEP...10..192C", "2017JHEP...11..134C", "2017JHEP...11..189K", "2017JHEP...11..207C", "2018JHEP...08..070K", "2018JHEP...12..042D", "2019JHEP...05..101C", "2019PhLB..79834997H", "2019arXiv190101454C", "2023arXiv231109937B", "2024PhR..1059....1C" ]
[ "astronomy", "physics" ]
3
[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]
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[ "10.48550/arXiv.1607.00004" ]
1607
1607.00004_arXiv.txt
\label{sec:Conclusions} In this work we focused on supersymmetric models where the LSP is a higgsino-like neutralino which plays the role of DM in the context of a non-standard cosmology. The difference with respect to the standard cosmological history comes from the presence of new degrees of freedom which can decay late changing the DM relic abundance produced by the standard thermal freeze-out scenario. The presence of such fields is well motivated from string theory where moduli fields naturally emerge in its low-energy 4D limit. The paper is divided into two parts. In Sec. \ref{section1}, \ref{section2} and \ref{sec:IndirectDetection} we performed a model independent analysis of supersymmetric models with non-thermal production of light higgsino DM. In Sec. \ref{sec:DirectDetection} we presented instead a model dependent discussion of different string models where a non-standard cosmology is motivated by the presence of moduli which decay at late time. For each string model we studied theoretical and observational constraints on higgsino non-thermal DM production. The main conclusions of the model independent analysis developed in the first part of the paper are: \begin{enumerate} \item In non-thermal cosmologies with an extra period of matter domination which ends via reheating with temperatures of $\mc{O}(1 - 10)$ GeV (above BBN), light higgsinos with masses as low as a few hundred GeV can correctly saturate the DM content measured by the Planck satellite. \item Such light higgsinos are very interesting from both a theoretical and an experimental point of view. The fact that they are very light makes them easily accessible to both indirect detection and collider searches. \item The strongest bound from indirect detection imposes that non-thermally produced higgsinos cannot be lighter than $300$ GeV. This bound comes from Fermi-LAT dSph where the dependence on the DM astrophysical profile is less important than in galactic centre observations. Observations by future experiments like CTA, together with data from Fermi-LAT GC, could cover essentially the entire parameter space of this scenario. On the other hand, unlike in the thermal case, collider signals from a $100$ TeV machine could test directly this scenario using searches on monojet and disappearing tracks. \end{enumerate} From the model dependent discussion performed in the second part of the paper, we can conclude that: \begin{enumerate} \item The main difference between LVS and KKLT scenarios for type IIB moduli stabilisation is that the last particle to decay in LVS models is the lightest modulus, while in KKLT models it is the gravitino. However, both cases feature a late decaying particle (scalar in LVS and fermion in KKLT) which motivates the analysis performed in the first part of the paper. Depending on the scenario under consideration, the hierarchy between the masses of the moduli, the higgsinos and the other superpartners can take a different form. \item When the visible sector is localised on D7-branes, both LVS and KKLT models with stable higgsino LSP are plagued by the problem of DM overproduction since heavy gaugino masses give rise a large contribution to higgsino masses at one-loop level. \item KKLT models with the visible sector on D3-branes still tend to have problems with higgsino DM overproduction due to the fact that gauginos are heavy in order to have gravitinos which decay before BBN. However, there is a fine-tuned region of the underlying parameter space where the non-thermal production of light higgsinos can yield the correct DM abundance. \item LVS models with the visible sector on D3-branes seem to be the best option to realise non-thermal scenarios with light higgsino DM. In fact, one-loop corrections to higgsino masses are small since sequestering effects suppress gaugino masses with respect to the mass of the decaying modulus. By exploiting the relation between the modulus and the gaugino mass, we managed to rewrite the reheating temperature in terms of the gaugino mass. This allowed us to introduce the effect of DM direct detection searches. We have found that, on the one hand, it is necessary to use large scale DM direct detection experiments (beyond $1$ Ton) to constrain more than what indirect detection already does, while, on the other hand, a large region of the parameter space falls below the neutrino background, and so DM direct detection experiments do not seem to be very useful to explore the parameter space of these theories. \end{enumerate} Future experiments will be able to completely probe the underlying parameter space of supersymmetric models with non-thermal light higgsino DM. This makes this scenario very interesting from both DM detection and future collider searches at $100$ TeV and motivates a detailed analysis from both sides.
16
7
1607.00004
We study the scenario of higgsino dark matter in the context of a non-standard cosmology with a period of matter-domination prior to Big-Bang nucleosynthesis. Matter-domination changes the dark matter relic abundance if it ends via reheating to a temperature below the higgsino thermal freeze-out temperature. We perform a model independent analysis of the higgsino dark matter production in such scenario. We show that light higgsino-type dark matter is possible for reheating temperatures close to 1 GeV. We study the impact of dark matter indirect detection and collider physics in this context. We show that Fermi-LAT data rules out non-thermal higgsinos with masses below 300 GeV. Future indirect dark matter searches from Fermi-LAT and CTA would be able to cover essentially the full parameter space. Contrary to the thermal case, collider signals from a 100 TeV collider could fully test the non-thermal higgsino. In the second part of the paper we discuss the motivation of such non-thermal cosmology from the perspective of string theory with late-time decaying moduli for both KKLT and LVS moduli stabilization mechanisms. We describe the impact of embedding dark matter higgsino in these scenarios.
false
[ "non-thermal higgsinos", "such non-thermal cosmology", "dark matter higgsino", "higgsino dark matter", "dark matter indirect detection", "LVS moduli stabilization mechanisms", "Future indirect dark matter searches", "temperatures", "the non-thermal higgsino", "such scenario", "light higgsino-type dark matter", "a non-standard cosmology", "the higgsino dark matter production", "the dark matter relic abundance", "string theory", "late-time decaying moduli", "KKLT", "Matter-domination", "matter-domination", "the thermal case" ]
8.584244
-1.888211
54
12516901
[ "Agapito, G.", "Puglisi, A.", "Esposito, S." ]
2016SPIE.9909E..7EA
[ "PASSATA: object oriented numerical simulation software for adaptive optics" ]
18
[ "Osservatorio Astrofisico di Arcetri (Italy)", "Osservatorio Astrofisico di Arcetri (Italy)", "Osservatorio Astrofisico di Arcetri (Italy)" ]
[ "2016SPIE.9909E..1BR", "2016SPIE.9909E..3UE", "2016SPIE.9909E..6BE", "2017MNRAS.472.2544P", "2019arXiv191105989A", "2020SPIE11448E..2SA", "2020SPIE11448E..3VT", "2021arXiv210107091P", "2022JATIS...8b1503S", "2022JATIS...8b1505A", "2022SPIE12185E..56A", "2022SPIE12185E..5VT", "2023A&A...677A.168A", "2023JATIS...9d9005A", "2023aoel.confE..10A", "2023aoel.confE..29C", "2023arXiv230200335B", "2024arXiv240615336A" ]
[ "astronomy", "physics" ]
6
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "2010SPIE.7736E..09E", "2010SPIE.7736E..3HQ", "2015MmSAI..86..446E", "2016SPIE.9909E..1BR", "2016SPIE.9909E..3UE", "2016SPIE.9909E..3VP" ]
[ "10.1117/12.2233963", "10.48550/arXiv.1607.07624" ]
1607
1607.07624_arXiv.txt
\label{sec:intro} PASSATA is an IDL based library/software capable of doing Monte-Carlo end-to-end Adaptive Optics (AO) simulations. PASSATA was originally developed to evaluate the performance of the Large Binocular Telescope (LBT) First Light Adaptive Optics (FLAO) system\cite{FLAO}. This system is a single conjugate adaptive optics system with a pyramid wavefront sensor. The first version was a library of functions called by a very long batch file with very limited flexibility. When the AO group of the Arcetri observatory started to work on other projects (for example: ARGOS\cite{rabienARGOS}, ERIS\cite{ERISSPIE2016}, GMT NGSAO\cite{doi:10.1117/12.2057059}), PASSATA was rewritten to be more flexible and more user friendly. The library core was kept the same, but it was wrapped in separated modules taking advantage of the possibilities of object-oriented programming made available by the last versions of IDL. In addition, the most computationally demanding routines were reimplemented using GPUs.
The PASSATA simulator is now a flexible and powerful tool, that can simulate most existing or in-development AO systems, and that can be easily extended with new features, thanks to the object-oriented approach and the very loose coupling between different parts of the code. GPU acceleration provides typical speedups of 50x to 100x w.r.t CPU code, and allows the study of the largest AO system currently in development. The current bottleneck for simulator performance is the speed and memory capacity of currently available GPUs, and special cases like EELT systems with non-point like sources stretch the simulation time to the practical limit. However GPU vendors have encourangingly robust develoment roadmaps for the near future, and thus we expect these special cases to benefit from future hardware releases.
16
7
1607.07624
We present the last version of the PyrAmid Simulator Software for Adaptive opTics Arcetri (PASSATA), an IDL and CUDA based object oriented software developed in the Adaptive Optics group of the Arcetri observatory for Monte-Carlo end-to-end adaptive optics simulations. The original aim of this software was to evaluate the performance of a single conjugate adaptive optics system for ground based telescope with a pyramid wavefront sensor. After some years of development, the current version of PASSATA is able to simulate several adaptive optics systems: single conjugate, multi conjugate and ground layer, with Shack Hartmann and Pyramid wavefront sensors. It can simulate from 8m to 40m class telescopes, with diffraction limited and resolved sources at finite or infinite distance from the pupil. The main advantages of this software are the versatility given by the object oriented approach and the speed given by the CUDA implementation of the most computational demanding routines. We describe the software with its last developments and present some examples of application.
false
[ "Pyramid wavefront sensors", "several adaptive optics systems", "ground based telescope", "a single conjugate adaptive optics system", "end", "Adaptive Optics", "Shack Hartmann", "a pyramid wavefront sensor", "CUDA", "Arcetri", "single conjugate, multi conjugate and ground layer", "Adaptive", "an IDL and CUDA based object oriented software", "PASSATA", "Pyramid", "the object oriented approach", "sources", "IDL", "application", "finite or infinite distance" ]
14.325491
10.557856
10
12508324
[ "Zhu, Zhen-Yu", "Li, Ang", "Hu, Jin-Niu", "Sagawa, Hiroyuki" ]
2016PhRvC..94d5803Z
[ "Δ (1232 ) effects in density-dependent relativistic Hartree-Fock theory and neutron stars" ]
58
[ "Department of Astronomy, Xiamen University, Xiamen 361005, China", "Department of Astronomy, Xiamen University, Xiamen 361005, China", "School of Physics, Nankai University, Tianjin 300071, China", "RIKEN Nishina Center, RIKEN, Wako 351-0198, Japan; Center for Mathematics and Physics, University of Aizu, Aizu-Wakamatsu, Fukushima 965-8560, Japan" ]
[ "2016arXiv161008770L", "2017ApJ...844...41L", "2017ApJ...850...20G", "2017arXiv170202310S", "2018ApJ...862...98Z", "2018EPJA...54..133L", "2018IAUS..337..360L", "2018PhLB..783..234L", "2018PhRvC..97f5202T", "2018PhRvC..98d5801S", "2018PhRvD..97b3018B", "2018arXiv180107084J", "2019AIPC.2127b0020S", "2019ApJ...877..139G", "2019ApJ...883..168R", "2019ApJ...883..174X", "2019IJMPD..2850040S", "2019IJMPD..2850122S", "2019PhRvC.100a5809L", "2019PhRvD..99b3004S", "2019PrPNP.10903714B", "2019arXiv190306057J", "2020ApJ...893...61Z", "2020ApJ...902...38Z", "2020JHEAp..28...19L", "2020MNRAS.499..914R", "2020PhLB..80635517L", "2020PhLB..81035812L", "2020PhRvD.102f3008M", "2021EPJA...57..309C", "2021JPhG...48j5201S", "2021MNRAS.507.2991T", "2021PhLB..81436070R", "2021PhRvC.103d5804S", "2021PhRvC.104e5805S", "2021PhRvD.103f3004T", "2021Univ....7..182L", "2021Univ....7..382S", "2021Univ....7..408L", "2021arXiv210208787B", "2021arXiv211212629B", "2022ApJ...925...16T", "2022ApJ...935...88H", "2022EPJA...58..115R", "2022EPJA...58..137S", "2022Galax..10...16K", "2022IJMPE..3150050K", "2022MPLA...3750218H", "2022PhRvC.105a5802T", "2022PhRvD.105j3009M", "2022arXiv220302272B", "2023MNRAS.522.3263I", "2023MNRAS.525.5512I", "2023PrPNP.13104041S", "2023Symm...15.1123K", "2024PhRvD.109b3004C", "2024PhRvD.109l3005M", "2024arXiv240214288W" ]
[ "astronomy", "physics" ]
6
[ "Nuclear Theory", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1974PhRvD...9.1613M", "1979NuPhA.322..382G", "1981PhLB..100...10B", "1987PhRvC..36..380B", "1989NuPhA.504..797W", "1995PhRvC..51.2260J", "1997PhRvC..55.2909W", "1997PhRvC..56.1570L", "1999NuPhA.656..331T", "2003PhRvC..67c8801X", "2005PhLB..615..193Y", "2006PhLB..639..242L", "2006PhLB..640..150L", "2006PhRvC..74e5801L", "2007PhRvC..76c4314L", "2008PhRvC..77f5807P", "2008PhRvC..78f5805S", "2008PhRvL.101l2502L", "2010ApJ...724L..74S", "2010Natur.467.1081D", "2010PhRvC..81b5806L", "2010PhRvD..82j1301O", "2011PhRvC..83b5804B", "2011PhRvC..84c5801S", "2012APh....37...70L", "2012PhRvC..85b5806L", "2012PhRvC..85f4606W", "2012PhRvC..86c5502Z", "2013ApJ...771...51L", "2013ApJ...772....7G", "2013ApJ...774...17S", "2013Sci...340..448A", "2014ApJ...784..123L", "2014PhRvC..89b5802H", "2014PhRvC..90f5809D", "2014PhRvD..89d3014D", "2015PhRvC..91c5803L", "2015PhRvC..92a5802C", "2015PhRvC..92c4603L", "2015PhRvL.114i2301L", "2016ARA&A..54..401O", "2016ApJ...832..167F", "2016EPJA...52...21R" ]
[ "10.1103/PhysRevC.94.045803", "10.48550/arXiv.1607.04007" ]
1607
1607.04007_arXiv.txt
Recently the early appearance of $\Delta$-isobars in dense nuclear matter has intrigued many illuminating studies for the interest of both neutron stars (NSs)~\cite{drago14prc,schramm10apj,drago14prd,guo03prc,li15prc,delta} and heavy ion collision~\cite{hic}. In particular, very compact stellar configurations~\cite{drago14prc,schramm10apj} are reached due to the introduction of $\Delta$-isobars as required by some current radius measurements~\cite{smallradius1,smallradius2,smallradius3}. However, there is the $\Delta$ puzzle~\cite{drago14prc} similar with the hyperon puzzle~\cite{hpuzzle,Li11y, rijken} that the corresponding maximum NSs' masses could not fulfill the constraint from the recent two precisely-measured 2-solar-mass pulsars PSR J1614-2230~\cite{2solar10,2solar11} and PSR J0348+0432~\cite{2solar2}. Previous studies have suggested the two-family scenario of compact stars to resolve this problem~\cite{drago14prd}, we here go beyond the previously-employed relativistic mean field (RMF) theory to include the contributions due to the exchange (Fock) terms and the pseudo-vector $\pi$-meson couplings which are effective only through exchange terms. Therefore, we employ the density-dependent relativistic Hartree-Fock (DDRHF) theory~\cite{bouyssy87prc} to study this problem. As one of the most advanced nuclear many-body model based on the covariant density functional theory, DDRHF presents a quantitative description of nuclear phenomena~\cite{long12prc,sun08prc,long06plb,long08el,ddrhf1,ddrhf2,ddrhf3} with a similar accuracy as RMF. In DDRHF, the Lorentz covariant structure is kept in full rigor, which guarantees all well-conserved relativistic symmetries. Previously, it has been demonstrated that the isoscalar Fock terms are essential for the prediction of NS properties~\cite{long12prc,sun08prc}. Also, we expect that the density-dependence~\cite{long06plb,rmf99npa} introduced in the present study for all meson couplings would make difference on the $\Delta$-matter study. Besides the $\Delta$-puzzle problem in NSs, we are also interested in how the DDRHF results depend on the uncertain $\Delta$-meson coupling constants, since there are hardly previous studies except with RMF. Presently, there is no constraint on $\Delta$-$\rho$ coupling. Several constraints exist for $\sigma, \omega$ couplings from the quark model~\cite{1prd74,1npa79,plb81x}, finite-density QCD sum-rule methods~\cite{qcd95}, quantum hadrodynamics~\cite{prc12x,npa89x}, and laboratory experiments~\cite{prc12xx,prc97x,plb05x}. Based on the quark counting argument~\cite{1prd74,1npa79}, there are the universal couplings between nucleons, $\Delta$-isobars with mesons, namely, $x_{\lambda} = g_{\lambda\Delta}/g_{\lambda N} = 1$ ($\lambda$ labels meson). A theoretical analysis~\cite{plb81x} of M1 giant resonance and Gamow-Teller transitions in nuclei found $25-40\%$ reductions of the transition strength due to the couplings to $\Delta$ isobars, while recent experiments~\cite{prc12xx,prc97x,plb05x} bring down the quenching value to be at most $10-15\%$ due to the coupling to $\Delta$-isobars. This means that the $\Delta$ couplings of isoscalar mesons ($\sigma, \omega$) should be weaker than those of quark model prediction, namely $x_{\sigma}, x_{\omega}\leq1$ are suggested by these experiments. Also, at the saturation density $\rho_0$, a possible smaller $\Delta$-isobars' vector self-energy was indicated~\cite{qcd95} than the corresponding value for the nucleon [i.e., $x_{\omega}(\rho_0)<\sim1$], while the $\Delta$-isobars' scalar self-energy was difficult to predict. In the analysis of Ref.~\cite{npa89x} in Hartree approximation, the difference between $x_{\sigma}$ and $x_{\omega}$ was found to be $x_{\omega}=x_{\sigma}-0.2$. A recent study, however, concluded that the cross sections involved are not sensitive to $x_{\omega}$ and $x_{\sigma}$ either~\cite{prc12x}. In the present study, for the purpose of comparing with previous RMF studies~\cite{li15prc}, we scale the density-dependence of $\Delta$-meson coupling in a reasonable range according to previous constraints for $\sigma, \omega, \rho$ mesons. The $\pi$ meson coupling is fixed to be $x_\pi=1.0$ hereafter. We provide the DDRHF formula extended to include $\Delta$-isobars in Sect.~II. Our results and discussions will be given in Sect.~III, before drawing conclusions in Sect.~IV.
Summarizing, we have extended the DDRHF theory to include the $\Delta$-isobars for the interest of dense matter and NSs. For this purpose, we solve the Rarita-Schwinger equation for spin-3/2 particle with the $\Delta$ self-energy determined self-consistently. Four mesons (the isoscalar $\sigma$ and $\omega$ as well as isovector $\rho$ and $\pi$) are included for producing interaction between the baryons. All four available parameter sets (PKO1, PKO2, PKO3, PKA1) are employed to explore the model utilization. Also, the calculations are done with various choices of the uncertain meson-$\Delta$ couplings following the constrains reported in the various analysis of experimental data. We found that $\Delta$-isobars appear early in the nuclear matter, soften the EoS of NS matter, and reduce the NSs' maximum mass. We observed a controlled behaviour of the results, with respect to the change of the model parameter set and different meson-$\Delta$ couplings. In particular, the $\Delta$ softening of the EoS is due to the large decrease in Fock channel, mainly from the isoscalar mesons. On the other hand, the pion contributions are found to be negligibly small. Finally, we conclude that within DDRHF, a NS with $\Delta$-isobars in its core can be as heavy as the two recent massive stars whose masses are accurately measured, with a slightly smaller radius than the corresponding normal NSs.
16
7
1607.04007
The density-dependent relativistic Hartree-Fock (DDRHF) theory is extended to include Δ isobars for the study of dense nuclear matter and neutron stars. To this end, we solve the Rarita-Schwinger equation for spin-3/2 particle. Both the direct and exchange terms of the Δ isobars' self-energies are evaluated in detail. In comparison with the relativistic mean field theory (Hartree approximation), a weaker parameter dependence is found for DDRHF. An early appearance of Δ isobars is recognized at ρ<SUB>B</SUB>∼0.28 fm<SUP>-3</SUP>, comparable with that of hyperons. Also, we find that the Δ isobars' softening of the equation of state is mainly due to the reduced Fock contributions from the coupling of the isoscalar mesons, while the pion contributions are negligibly small. We finally conclude that with typical parameter sets, neutron stars with Δ isobars in their interiors could be as heavy as the two massive pulsars whose masses are precisely measured, with slightly smaller radii than normal neutron stars.
false
[ "normal neutron stars", "neutron stars", "Δ isobars", "dense nuclear matter", "hyperons", "detail", "spin-3/2 particle", "Hartree approximation", "typical parameter sets", "DDRHF", "slightly smaller radii", "the reduced Fock contributions", "SUP>-3</SUP", "fm", "the Δ isobars self-energies", "the pion contributions", "Hartree", "state", "the Δ isobars softening", "the isoscalar mesons" ]
4.635037
1.930073
71
12406460
[ "Louvet, F.", "Dougados, C.", "Cabrit, S.", "Hales, A.", "Pinte, C.", "Ménard, F.", "Bacciotti, F.", "Coffey, D.", "Mardones, D.", "Bronfman, L.", "Gueth, F." ]
2016A&A...596A..88L
[ "ALMA observations of the <ASTROBJ>Th 28</ASTROBJ> protostellar disk. A new example of counter-rotation between disk and optical jet" ]
22
[ "Departamento de Astronomia de ChileUniversidad de Chile, Santiago, Chile", "UMI-FCA, CNRS/INSU, UMI, 3386, France; Departamento de Astronomia de ChileUniversidad de Chile, Santiago, Chile; Univ. Grenoble Alpes, CNRS, IPAG, 38000, Grenoble, France", "Laboratoire d'Études du Rayonnement et de la Matière en Astrophysique et Atmosphères (LERMA), Observatoire de Paris-Meudon, Paris, France; Univ. Grenoble Alpes, CNRS, IPAG, 38000, Grenoble, France", "Atacama Large Millimeter/Submillimeter Array, Joint ALMA Observatory, Alonso de Córdova 3107, Vitacura 763-0355, Santiago, Chile; National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, Virginia, 22903-2475, USA", "UMI-FCA, CNRS/INSU, UMI, 3386, France; Departamento de Astronomia de ChileUniversidad de Chile, Santiago, Chile; Univ. Grenoble Alpes, CNRS, IPAG, 38000, Grenoble, France", "UMI-FCA, CNRS/INSU, UMI, 3386, France; Departamento de Astronomia de ChileUniversidad de Chile, Santiago, Chile; Univ. Grenoble Alpes, CNRS, IPAG, 38000, Grenoble, France", "Instituto Nazionale di Astrofisica - Osservatorio Astrofisico di Arcetri, 50125, Firenze, Italy", "School of Physics, University College Dublin, Belfield, Dublin 2, Ireland", "Departamento de Astronomia de ChileUniversidad de Chile, Santiago, Chile", "Departamento de Astronomia de ChileUniversidad de Chile, Santiago, Chile", "Institut de Radioastronomie Millimétrique-Grenoble, 38400 Saint-Martin-d', Hères, France" ]
[ "2016A&A...596A..92K", "2016ApJ...832..152D", "2017A&A...607L...6T", "2017MmSAI..88..801M", "2018A&A...618A.120L", "2018ApJ...868...28F", "2018sf2a.conf..311L", "2019ApJ...870...76B", "2020A&A...640A..82T", "2020A&A...640A.111A", "2020ApJ...903...78P", "2020MNRAS.494..827M", "2021A&A...648A..41V", "2021A&A...652A.119M", "2021ApJ...909..196L", "2021NewAR..9301615R", "2022A&A...658A.102O", "2022MNRAS.514.1088Z", "2023A&A...669A..81O", "2023A&A...673A.156M", "2023A&A...677A.129O", "2023ASPC..534..567P" ]
[ "astronomy" ]
11
[ "stars: individual: Th 28", "ISM: jets and outflows", "techniques: interferometric", "submillimeter: ISM", "stars: formation", "circumstellar matter", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1977ApJS...35..161S", "1982MNRAS.199..883B", "1983QJRAS..24..267H", "1986A&A...161..195K", "1988PASP..100.1529G", "1991ApJ...381..250B", "1993ApJ...402..280B", "1994A&A...291L..23G", "1994AJ....108.1071H", "2000A&A...355..165D", "2000A&A...358..593S", "2001A&A...369..993P", "2002A&A...394L..31T", "2002ApJ...576..222B", "2003A&A...399..773D", "2003ApJ...590L.107A", "2004ApJ...604..758C", "2005A&A...432..149W", "2005ApJ...631.1134A", "2006A&A...446..971J", "2006A&A...448..231C", "2006A&A...452..897C", "2006A&A...453..785F", "2006ASPC..351..319M", "2007ApJ...663..350C", "2007prpl.conf..231R", "2007prpl.conf..261S", "2007prpl.conf..277P", "2008ApJS..177..551M", "2009AJ....138.1072W", "2009ASSP...13..247C", "2009ApJ...701..260I", "2009MNRAS.399.1802R", "2010A&A...511A..10C", "2011IAUS..275..383F", "2011MNRAS.418.1194M", "2012ApJ...749..139C", "2012ApJ...759L...1S", "2013A&A...558A..77G", "2014MNRAS.442...28W", "2015ApJ...804....2C", "2015MNRAS.446.3975S" ]
[ "10.1051/0004-6361/201628474", "10.48550/arXiv.1607.08645" ]
1607
1607.08645_arXiv.txt
\label{s:intro} \begin{figure*}[htb!] \subfloat{\includegraphics[scale=0.32]{12co.pdf}} \hspace{0.3cm}\subfloat{\includegraphics[scale=0.32]{13co.pdf}} \hspace{0.3cm}\subfloat{\includegraphics[scale=0.32]{continuum.pdf}} \caption{\textbf{Left:} moment zero of the $^{12}$CO(3-2) integrated from -15 to 20 km.s$^{-1}$. Contours start at 5\,$\sigma$ with 20\,$\sigma$ steps. \textbf{Middle:} moment zero of the $^{13}$CO(3-2) integrated from -15 to 20 km.s$^{-1}$. Contours start at 5\,$\sigma$ with 5\,$\sigma$ steps. \textbf{Right:} continuum map. Contours start at 5\,$\sigma$ with 20\,$\sigma$ steps. Noise level at 1\,$\sigma$ is indicated in the upper right corner of each panel; the unit is mJy/beam.km/s in the two left panels and mJy/beam in the right panel. The black ellipses on the three panels represent the FWHM intensity contours. White hatched ellipses represent the clean beam FWHM.} \label{f:the3} \end{figure*} The interplay between accretion and ejection of matter is believed to be a crucial element in the formation of stars. Yet, the exact link between jets and accretion disks is still a critical issue in contemporary astrophysics \citep{ray07}. The most attractive possibility is a fundamental connection via angular momentum transfer from disk to jet by means of magnetocentrifugal forces, so that circumstellar material may continue to accrete onto the central object \citep[e.g.,][]{blandford82}. Exactly where this transfer occurs and how it impacts the disk physics is, however, still hotly debated \citep{ferreira06, pudritz07, shang07, romanova09, cabrit09}. Owing to their low extinction and their proximity ($\approx$ 150 pc), accreting classical T Tauri stars (hereafter CTTSs) offer a unique laboratory to elucidate the connection between accretion and ejection. A recent major breakthrough was the detection with the STIS instrument mounted on the HST of transverse Doppler shifts of 5 - 30\, km.s$^{-1}$ in optical and near-ultraviolet (hereafter NUV) emission lines on scales of 50 AU across the base of four close CTTS jets \citep{bacciotti02, woitas05, coffey04, coffey07}. If these shifts are due to jet rotation in a steady state outflow, they would imply that jets are magnetically launched from the disk surface at relatively small-scale radii of 0.1-1.6 AU, thus solving the long-standing problem of the jet origin in CTTS \citep{bacciotti02,anderson03,coffey04,woitas05}. However, effects other than rotation \citep[e.g., precession; see][]{cerqueira06} could cause transverse Doppler shifts similar to those detected in T Tauri jets. It is crucial that the underlying disks rotate in the same sense as the jet to support the jet rotation interpretation. This comparison was conducted in three of the four systems mentioned above (\object{CW\,Tau}, \object{DG\,Tau}, and \object{RW\,Aur}), plus in \object{RY\,Tau}. In \object{RY Tau}, the study was inconclusive as no systematic Doppler shifts were found in the jet; see \cite{coffey15}. The sense of rotation of the jets of \object{CW\,Tau} and \object{DG\,Tau} \citep{coffey07}, traced by optical lines, coincides with the sense of rotation of their disks (\citealt{testi02}; Cabrit private comm., respectively). The jet rotation sense traced in near ultra-violet (NUV) emission lines at higher flow velocities is coherent for \object{DG\,Tau}, but these emission lines are below the detection threshold for \object{CW\,Tau} \citep{coffey07}. In RW\,Aur, the optical jet and disk rotations appear opposite (see \citealt{coffey04} and \citealt{cabrit06}). Here, jet rotation traced in the NUV was found to be compatible with the disk rotation sense (i.e., opposite to the optical jet rotation), but could only be measured in one lobe and at one of the two epochs \citep[spaced by six months; see ][]{coffey12}; because of its complexity and variability the authors declared \object{RW\,Aur} as unsuitable for jet rotation studies. We present the confrontation of jet and disk rotation senses for the fourth system, \object{Th\,28}. The T~Tauri star \object{ThA 15-28} \citep[hereafter \object{Th\,28};][]{the62}, also known as \object{Sz~102} \citep{schwartz77} or Krautter's star, is a young member of the \object{Lupus~3} association. Although uncertain, as the photospheric spectrum is veiled, a spectral type G8 to K2 was quoted for the central object of \object{Th\,28} \citep{graham88,hughes94,mortier11}. Even for a K star, the stellar emission from \object{Th\,28} is significantly underluminous: 0.03 L$_{\odot}$ \citep{mortier11}, which suggests that the disk is observed close to edge-on and obscures the central star. The derived disk luminosity, inferred from the spectral energy distribution (SED) integrated excess above the stellar photosphere, is $\sim$0.13 L$_{\odot}$ \citep{mortier11, merin08}. We adopt the recent value of d=185$^{\rm +11}_{\rm -10}$\,pc derived by \citet{Galli13} for the Lupus~3 cloud, with an improved convergent point search method, for the distance to the source. \object{Th\,28} drives a bright and striking bipolar jet whose axis lies at PA=98$^{\circ}$ \citep{krautter86,graham88}. The large proper motions derived for the distant knots are compatible with a close to edge-on geometry \citep[e.g.,][]{comeron10}. \object{Th\,28} is considered one of the best case for jet rotation detection in the HST sample, as transverse velocity gradients are well detected in both lobes of its optical atomic jet and are consistent in different optical lines \citep{coffey04}. In optical, the jet rotation sense is therefore well determined with a clockwise direction looking down the blueshifted jet lobe toward the source. The situation for NUV lines were inconclusive \citep{coffey07}. Only one measurement on the receding jet, that is, however, consistent with the sense of rotation derived in optical, exceeded the incertitude level. \smallskip In this paper, we report ALMA band 7 CO and continuum observations at 0.4$^{\prime\prime}$-0.5$^{\prime\prime}$ angular resolution of the \object{Th\,28} system, aimed to detect the disk and determine its sense of rotation. We detail our observations and data reduction in Sect.~\ref{s:obs}, and demonstrate that the CO and continuum emissions clearly trace a disk in Keplerian rotation around \object{Th\,28} in Sect.~\ref{s:results}. The Sect.~\ref{s:analysis} offers a detailed analysis of the disk properties. In Sect.~\ref{s:discu} we discuss the implications for jet launching models and central source properties. We summarize our conclusions in Sect.~\ref{s:concl}.
\label{s:concl} We have observed the T Tauri star \object{Th\,28} during the ALMA-cycle 1 campaign in Band 7 (275-373\,GHz). We detected $^{12}$CO($J$=3$\rightarrow$2), $^{13}$CO($J$=3$\rightarrow$2) and continuum signatures of a Keplerian accretion disk around \object{Th\,28}. \begin{itemize} \item[$\bullet$] The $^{12}$CO emission is clearly resolved and elongated along PA=7.3$^{\circ}$. The morphology of the disk seen in CO matches very well the morphology derived for the large-scale atomic jet. The PA of the disk is perpendicular to that of the jet, and both inclinations are comparable. \item[$\bullet$] The $^{12}$CO shows kinematics that are consistent with a disk in Keplerian rotation. Indeed its double-peak integrated profile, plus its PV diagram along the PA of the disk, are characteristic of Keplerian rotation. \item[$\bullet$] We derive a V$_{\rm lsr}$=3$\pm$0.1 km.s$^{-1}$ for the central source in agreement with the range of values observed in the Lupus~3 cloud. \item[$\bullet$] We constrain the power law index, $q$, of the temperature distribution to $q\simeq$0.8 and the R$_{\rm out}$ to $\simeq$ 95\,AU. \item[$\bullet$]The combination of large R$_{\rm out}$ and peak velocity separation in $^{12}$CO of $\Delta$v=7\,km.s$^{-1}$ is best reproduced with a central stellar mass of 1.4-1.8\,M$_\odot$. This is also consistent with early K spectral type estimates for this source. \item[$\bullet$]The rotation sense of the disk is well detected with our ALMA CO observations and this direction is opposite that of the transverse velocity shifts previously detected with HST in the optical jet. This discrepancy in rotation senses implies that velocity gradients measured in optical lines cannot be used to infer launching radii. The NUV domain, which probes a faster and more collimated inner part of the jet, is likely more reliable to measure jet rotation from Doppler gradients. \end{itemize}
16
7
1607.08645
<BR /> Aims: Recently, differences in Doppler shifts across the base of four close classical T Tauri star jets have been detected with the HST in optical and near-ultraviolet (NUV) emission lines, and these Doppler shifts were interpreted as rotation signatures under the assumption of steady state flow. To support this interpretation, it is necessary that the underlying disks rotate in the same sense. Agreement between disk rotation and jet rotation determined from optical lines has been verified in two cases and rejected in one case. Meanwhile, the near-ultraviolet lines, which may trace faster and more collimated inner spines of the jet than optical lines, either agree or show no clear indication. We propose to perform this test on the fourth system, <ASTROBJ>Th 28</ASTROBJ>. <BR /> Methods: We present ALMA high angular resolution Band 7 continuum, <SUP>12</SUP>CO(3-2) and <SUP>13</SUP>CO(2-1) observations of the circumstellar disk around the T Tauri star <ASTROBJ>Th 28</ASTROBJ>. <BR /> Results: The sub-arcsecond angular resolution (0.46''× 0.37'') and high sensitivity reached enable us to detect, in CO and continuum, clear signatures of a disk in Keplerian rotation around <ASTROBJ>Th 28</ASTROBJ>. The <SUP>12</SUP>CO emission is clearly resolved, allowing us to derive estimates of disk position angle and inclination. The large velocity separation of the peaks in <SUP>12</SUP>CO, combined with the resolved extent of the emission, indicate a central stellar mass in the range 1-2 M<SUB>⊙</SUB>. The rotation sense of the disk is well detected in both <SUP>13</SUP>CO and <SUP>12</SUP>CO emission lines, and this direction is opposite to that implied by the transverse Doppler shifts measured in the optical lines of the jet. <BR /> Conclusions: The <ASTROBJ>Th 28</ASTROBJ> system is now the second system, among the four investigated so far, where counter-rotation between the disk and the optical jet is detected. These findings imply either that optical transverse velocity gradients detected with HST do not trace jet rotation or that modeling the flow with the steady assumption is not valid. In both cases jet rotation studies that rely solely on optical lines are not suitable to derive the launching radius of the jet.
false
[ "jet rotation", "jet rotation studies", "disk rotation", "optical lines", "rotation signatures", "steady state flow", "Keplerian rotation", "rotation", "optical transverse velocity gradients", "disk position angle", "Th", "four close classical T Tauri star jets", "the optical jet", "CO", "Doppler shifts", "the optical lines", "<", "ASTROBJ", "ALMA high angular resolution", "HST" ]
10.195439
11.727539
-1
2827080
[ "Martínez-Sykora, Juan", "De Pontieu, Bart", "Carlsson, Mats", "Hansteen, Viggo" ]
2016ApJ...831L...1M
[ "On the Misalignment between Chromospheric Features and the Magnetic Field on the Sun" ]
36
[ "Bay Area Environmental Research Institute, Sonoma, CA 94952, USA ; Lockheed Martin Solar and Astrophysics Laboratory, Palo Alto, CA 94304, USA;", "Lockheed Martin Solar and Astrophysics Laboratory, Palo Alto, CA 94304, USA ; Institute of Theoretical Astrophysics, University of Oslo, P.O. Box 1029 Blindern, NO-0315 Oslo, Norway", "Institute of Theoretical Astrophysics, University of Oslo, P.O. Box 1029 Blindern, NO-0315 Oslo, Norway", "Lockheed Martin Solar and Astrophysics Laboratory, Palo Alto, CA 94304, USA ; Institute of Theoretical Astrophysics, University of Oslo, P.O. Box 1029 Blindern, NO-0315 Oslo, Norway" ]
[ "2017A&A...599A.133A", "2017ApJ...845L..18D", "2017ApJ...847...36M", "2017ApJ...847..141S", "2017IAUS..327....1R", "2017Sci...356.1269M", "2018A&A...609A..18P", "2018A&A...618A..87K", "2018ApJ...857...73C", "2019A&A...621A...1R", "2019A&A...627A..25P", "2019A&A...630A..79P", "2019A&A...631A..33B", "2019arXiv191208650S", "2020A&A...633A..66N", "2020A&A...635L...2B", "2020A&A...638A..79N", "2020A&A...642A.220G", "2020ApJ...889...95M", "2020ApJ...898...32K", "2020arXiv200208441B", "2021ApJ...911...41G", "2021ApJ...921...39A", "2021JGRA..12629097S", "2021MNRAS.501.1940A", "2021SoPh..296...70R", "2021SoPh..296...84D", "2022A&A...662A..88V", "2022ARA&A..60..415T", "2022ApJ...929...88M", "2022PhDT.........3P", "2023A&A...672A..89K", "2023SSRv..219....1T", "2024A&A...681A.114G", "2024RSPTA.38230229H", "2024arXiv240517095C" ]
[ "astronomy" ]
6
[ "magnetohydrodynamics: MHD", "methods: numerical", "Sun: atmosphere", "Sun: magnetic fields", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1965RvPP....1..205B", "2002ApJ...572L.113G", "2003SPIE.4853..341S", "2006A&A...447.1111S", "2006A&A...450..805L", "2006SoPh..235...55S", "2006SoPh..236..415C", "2007A&A...473..625L", "2007ApJ...660L.169R", "2007cemf.book.....P", "2008A&A...480..515C", "2008SoPh..247..249W", "2008SoPh..247..269M", "2009ApJ...705.1183A", "2010A&A...517A..49H", "2010ApJ...718.1070H", "2010SoPh..261..363J", "2011A&A...527L...8D", "2011A&A...531A.154G", "2011ApJ...739...67J", "2012A&A...539A..39C", "2012ApJ...750....6C", "2012ApJ...753..161M", "2012SoPh..275....3P", "2012SoPh..275...17L", "2012SoPh..275..207S", "2013A&A...554A..22V", "2013SoPh..288..481R", "2014ApJ...784...30G", "2014ApJ...791...60K", "2014SoPh..289.2733D", "2015ApJ...811..106H", "2015RSPTA.37350055P", "2016A&A...585A...4C", "2016ApJ...826...51Z", "2016ApJ...826...61A", "2016ApJS..224...25A" ]
[ "10.3847/2041-8205/831/1/L1", "10.48550/arXiv.1607.02551" ]
1607
1607.02551_arXiv.txt
Optically thick chromospheric spectral lines such as \ion{Ca}{2} 8542\AA\ or H$\alpha$ are formed over a wide range of heights from the photospheric line wings to the middle or upper chromosphere line core. Observations in these lines show a dramatic transition from wing features that appear to be dominated by convective motions or acoustic waves in the high plasma $\beta$ regime (gas pressure is larger than magnetic pressure), to more linear features in the core of the lines that appear to trace magnetic field lines \citep[e.g.][]{luc2007,Cauzzi:2008jk}. Because these linear features are most often assumed to reveal the direction of the magnetic field, chromospheric structuring is increasingly being used to help constrain magnetic field extrapolation codes \citep[e.g.][]{Wiegelmann:2008ud,Jing:2011qa,Aschwanden:2016fj,Aschwanden:2016qy,Zhu:2016fk}. Such codes typically use non-linear force-free field extrapolation methods based on photospheric magnetic field measurements. Since the boundary conditions are most readily measured in the photosphere but the field is not necessarily of a force-free nature at those heights, various methods are used to preprocess the magnetic field measurements \citep[e.g.][]{Regnier:2013nx} or to incorporate magnetic field information from a more force-free boundary region such as the upper chromosphere \citep{Metcalf:2008bd}. The latter method is entirely dependent on the assumption that chromospheric features such as fibrils and spicules, which dominate the upper chromosphere, are well aligned with the magnetic field. For example, \citet{Aschwanden:2016qy} extrapolated magnetic field lines using photospheric vector field measurements from the Helioseismic Magnetic Imager \citep[HMI,][]{Scherrer:2012qf} onboard of the Solar Dynamic Observatory \citep[SDO,][]{Pesnell:2012nr} using the Vertical-Current Approximation Non-linear Force Free Field code \citep[VAC-NLFFF code,][]{Aschwanden:2016fj} which finds the best match of the extrapolated field lines with chromospheric and coronal features observed with the following imaging instruments: the Interferometric Bidimensional Spectrometre \citep[IBIS,][]{Cavallini:2006yq}, the Rapid Oscillation in the Solar Atmosphere instrument \citep[ROSA,][]{Jess:2010rt}, the Interface Region Imaging Spectrograph \citep[IRIS,][]{De-Pontieu:2014vn}, and Atmospheric Imager Assembly \citep[AIA,][]{Lemen:2012uq} onboard of SDO. Are these chromospheric structures really aligned with magnetic field lines? In the single-fluid MHD framework, the assumption that the upper chromosphere in the vicinity of network or plage is in the low plasma $\beta$ regime and the observed features thus aligned with the magnetic field, seems reasonable. However, there are observational clues that this may not always be the case. For example, \citet{de-la-Cruz-Rodriguez:2011qd} show that in some cases chromospheric fibrils do not necessarily follow the magnetic field structures. They used Stokes profile observations of \ion{Ca}{2} 8542~\AA\ from the Spectro-Polarimeter for INfrared and Optical Regions \citep[SPINOR,][]{Socas-Navarro:2006ul} at the Dunn Solar Telescope and CRisp Imaging Spectro-Polarimeter \citep[CRISP,][]{Scharmer:2006gf} in full Stokes mode at the Swedish 1-m Solar Telescope \cite[SST,][]{Scharmer:2003ve} and found that the misalignment of some fibrils with the magnetic field lines can be larger than 45 degrees. What causes such a large deviation from the magnetic field direction? It is already known that, in principle, the neutral population in the partially ionized chromosphere can have an impact on the force-free nature of the chromospheric field \citep{Arber:2009ve}. Is it possible that neutrals can also lead to misalignment of chromospheric features and the magnetic field? In order to better understand this misalignment of the magnetic field lines with upper chromospheric features, we performed 2D advanced radiative MHD simulations using the {\it Bifrost} code \citep{Gudiksen:2011qy}. We included the effects of the interaction between ions and neutrals in the magnetized and partially ionized gas of the middle to upper chromosphere by including ambipolar diffusion in the induction equation (the so-called generalized Ohm's law) of our MHD code. Our simulations show that while neutrals are mostly coupled to the magnetic field through collisions with ions, under certain conditions the collisional frequency between neutrals and ions is low enough that the ions can become somewhat decoupled from the magnetic field, thereby allowing the magnetic field to diffuse and magnetic energy to be dissipated into thermal energy \citep[see][among others]{cowling1957,Braginskii:1965ul,Parker:2007lr}. We find that our simulation naturally produces misalignment of the magnetic field and spicules. Our model thus provides an explanation for why some of the observations show field lines that are misaligned with chromospheric features, and which conditions can lead to such misalignment.
\label{sec:con} Our 2.5D radiative MHD simulation includes ion-neutral interaction effects and produces some examples of chromospheric features that are decoupled from the magnetic field direction. This is a result of ion-neutral interaction effects in the chromosphere and can occur when the ambipolar diffusion and the current perpendicular to the magnetic field lines are large, and the thermo-dynamic timescales are at least of the same order as the ambipolar velocity timescales. The simulated features that become misaligned from the magnetic field have lifetimes of order a few minutes. The time-scale becomes shorter in regions with large currents perpendicular to the magnetic field lines and large ambipolar diffusion due to low values of the temperature, ion-neutral collision frequency and/or ionization degree. Under these conditions the magnetic field may undergo evolution that is different from that of the thermodynamic structures. For example, decaying spicules may change their connectivity with ``different'' field lines crossing the spicule. In such a case, instead of following a space-time parabolic profile along a fixed direction, the spicule may show a horizontal displacement at the same time as they disappear (towards the end of their lifetime). This process can provide a natural explanation for the observations of \citet{de-la-Cruz-Rodriguez:2011qd} where fibrils or spicule-like features do not necessarily follow the magnetic field direction. Our model shows that the misalignment is not uniform in space or time, with some structures less affected. In addition, dynamic features appear to be typically more misaligned towards the end of their lifetime. This occurs in particular in regions with enough current perpendicular to the magnetic field, i.e., where there is large magnetic tension. Such conditions can also be expected in active regions with strong currents, such as in newly emerging active regions. However, it is a priori not clear how one can determine from the observations alone which features are likely not well aligned with the magnetic field. Future work will be needed to investigate how misalignment can be estimated based on observational clues, such as the temporal behavior of the structures or the presence of currents. Our results suggest that ion-neutral interaction effects may have a significant impact on magnetic field extrapolation methods. Within the chromosphere, the ambipolar diffusion shows strong variations in both space and time. In regions where the ambipolar diffusion is strong the magnetic field will be more potential. However, at the boundaries between regions of strong and weak ambipolar diffusion, the magnetic field lines may have strong changes in the connectivity. These variations in space and time of the ambipolar diffusion impact the magneto-thermodynamic processes in the chromosphere \citep{Martinez-Sykora:2016a}. % This was missing in previous numerical models \citep{Carlsson:2016rt}. It is unclear how such spatial and temporal complexity can be captured in field extrapolation methods that are based on photospheric vector magnetograms. In addition, magnetic field extrapolation codes that attempt to use the direction of chromospheric features in order to find the best match with the field lines may provide field configurations that are incorrect. This is because the magnetic field may not be well aligned with the chromospheric features due to the presence of ambipolar diffusion. As a result of this, the measurements of the misalignment between features and magnetic field extrapolation using these methods may be incorrect and in general underestimated \citep{Aschwanden:2016qy}. Future studies of these extrapolation codes should investigate this issue further by trying the method on synthetic data from this type of radiative MHD models that include ion-neutral interaction effects and comparing with the actual magnetic field. Many studies of the solar atmosphere, in particular the corona, are undertaken using 1D hydrodynamic loop models \citep[e.g.][]{Klimchuk:2014fk}. Such models are based on the assumption that the thermodynamic evolution occurs along magnetic field lines or tubes (1D). However, our results indicate that decoupling of the field lines from plasma is a frequent occurrence in the chromosphere. This will change the connectivity of each element along the 1D models. In addition, ambipolar diffusion fundamentally undermines the assumption that the plasma is tied to the field on timescales of many minutes. As we have shown, this is not always the case. Therefore, these 1D models cannot capture or mimic the physics of the processes that connect the corona to the chromosphere and transition region. \cite{Leake:2006kx} performed 2D simulations of flux emergence with ambipolar diffusion. They also notice that the magnetic field is more potential in the atmosphere due to the ambipolar dissipation. However, they get several orders of magnitude smaller currents for the case with ambipolar diffusion than without ambipolar diffusion whereas we get from 1.5 to 4 times smaller in the case for ambipolar diffusion than without ambipolar diffusion. This is most likely due to the highly simplified setup of the ambipolar diffusion in their model chromosphere and also due to their missing many of the chromospheric processes driven by the convective motion. Our results also have a potential impact on more advanced 3D radiative MHD models. This is because the magnetic field energy deposition in the corona in radiative MHD models will change as soon as ion-neutral interaction effects are introduced. For example, \citet{Peter:2015zv} noticed that larger domains show stronger flows and Doppler shifts for transition region EUV profiles (similar to observations) compared to smaller simulated domains \citep{Hansteen:2010uq,Hansteen:2015qv}. Ion-neutral interaction effects will impact these results and deeper investigations that include ion-neutral effects must be performed for computational domains of various sizes. Our simulation suffers from several limitations that need to be addressed. Our simulation does not include time dependent ionization which will impact the spatial and temporal distribution of ambipolar diffusion in the chromosphere \citep{Leenaarts:2007sf,Golding:2014fk}. The process described above is also strongly constrained to the two dimensions of the model, therefore the expansion of these models into three dimensions is needed. Finally, the Generalized Ohm's law is valid as long as the time scales are much larger than the ion-neutral collision frequencies, but in the transition region this may not always be fulfilled and ions may decouple from neutrals \citep{Martinez-Sykora:2012uq}. This can potentially alter these results.
16
7
1607.02551
Observations of the upper chromosphere show an enormous amount of intricate fine structure. Much of this comes in the form of linear features, which are most often assumed to be well aligned with the direction of the magnetic field in the low plasma β regime that is thought to dominate the upper chromosphere. We use advanced radiative magnetohydrodynamic simulations, including the effects of ion-neutral interactions (using the generalized Ohm’s law) in the partially ionized chromosphere, to show that the magnetic field is often not well aligned with chromospheric features. This occurs where the ambipolar diffusion is large, I.e., ions and neutral populations decouple as the ion-neutral collision frequency drops, allowing the field to slip through the neutral population; where currents perpendicular to the field are strong; and where thermodynamic timescales are longer than or similar to those of ambipolar diffusion. We find this often happens in dynamic spicule or fibril-like features at the top of the chromosphere. This has important consequences for field extrapolation methods, which increasingly use such upper chromospheric features to help constrain the chromospheric magnetic field: our results invalidate the underlying assumption that these features are aligned with the field. In addition, our results cast doubt on results from 1D hydrodynamic models, which assume that plasma remains on the same field lines. Finally, our simulations show that ambipolar diffusion significantly alters the amount of free energy available in the coronal part of our simulated volume, which is likely to have consequences for studies of flare initiation.
false
[ "field extrapolation methods", "chromospheric features", "such upper chromospheric features", "features", "linear features", "ambipolar diffusion", "intricate fine structure", "the chromospheric magnetic field", "the magnetic field", "the same field lines", "flare initiation", "thermodynamic timescales", "free energy", "results", "plasma", "1D hydrodynamic models", "ion-neutral interactions", "important consequences", "the field", "the upper chromosphere" ]
12.220616
15.246511
2
12473340
[ "Kalus, B.", "Percival, W. J.", "Bacon, D. J.", "Samushia, L." ]
2016MNRAS.463..467K
[ "Unbiased contaminant removal for 3D galaxy power spectrum measurements" ]
11
[ "Institute of Cosmology &amp; Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Institute of Cosmology &amp; Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Institute of Cosmology &amp; Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK", "Institute of Cosmology &amp; Gravitation, University of Portsmouth, Dennis Sciama Building, Portsmouth PO1 3FX, UK; Department of Physics, Kansas State University, 116 Cardwell Hall, Manhattan, KS 66506, USA; National Abastumani Astrophysical Observatory, Ilia State University, 2A Kazbegi Ave., GE-1060 Tbilisi, Georgia" ]
[ "2017MNRAS.465.1847E", "2018PhRvD..97l3540S", "2018arXiv181004263P", "2019JCAP...04..023M", "2019MNRAS.482..453K", "2020MNRAS.495.1613R", "2020MNRAS.499.5527T", "2021JCAP...11..027B", "2021MNRAS.503.5061W", "2021MNRAS.506.3439R", "2023MNRAS.524.2463B" ]
[ "astronomy" ]
5
[ "methods: statistical", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1992ApJ...398..169R", "1994ApJ...426...23F", "1998ApJ...499..555T", "2000ApJ...538..473L", "2003ApJS..148...97B", "2004PhRvD..69l3003S", "2007CSE.....9...90H", "2011arXiv1106.1706S", "2011arXiv1110.3193L", "2012ApJ...761...14H", "2012MNRAS.424..564R", "2013arXiv1308.0847L", "2014MNRAS.444....2L", "2014PhRvL.113v1301L", "2015MNRAS.453L..11B", "2015PhRvD..92h3532S", "2016MNRAS.456.2095E" ]
[ "10.1093/mnras/stw2008", "10.48550/arXiv.1607.02417" ]
1607
1607.02417_arXiv.txt
Galaxy surveys provide a rich store of information about the nature of the Universe, allowing us to constrain cosmological models with baryon acoustic oscillations (BAO), gravitational models with redshift space distortions (RSD) and inflationary models with primordial non-Gaussianity. A basic statistic containing large-scale structure information is the galaxy power spectrum $P(k)$, which is the 2-point function of the Fourier transformed density field. Future large-scale structure surveys, such as the Dark Energy Spectroscopic Instrument survey \citep[DESI]{Schlegel:2011zz,Levi:2013gra}, Euclid \citep{Laureijs:2011gra}\footnote{www.euclid-ec.org} and the Square Kilometre Array (SKA) \footnote{www.skatelescope.org}, will probe larger volumes, therefore allowing us to measure more Fourier modes of the galaxy density field. The observed galaxy field can be contaminated with fluctuations of non-cosmological origin, such as variations due to the galactic extinction and the stellar density. Often the contaminants are not known exactly (e.g. we may know the shape of the spurious mode but may not know its exact amplitude) which makes their exact removal impossible. These modes have the potential to strongly bias cosmological constraints derived from the clustering measurements, so we need to correct or suppress these misleading modes in a responsible way. We now introduce the basic mathematical problem that we wish to solve and introduce the main methods of removing contaminants discussed in literature. We assume that we have measured the galaxy density field as real numbers in configuration space, which we (fast) Fourier transform to obtain a Hermitian density field $F(\k)$. Furthermore, we assume that the contamination can be described by another Hermitian field $f(\k)$, such that the true density field is given by \begin{equation} D(\k)=F(\k)-\revision{\varepsilon_\mathrm{true}}f(\k), \label{eq:Fassum} \end{equation} \revision{with $\varepsilon_\mathrm{true}$ unknown.} In cases with multiple contaminants (which we label with capital Latin indices), we extend Eq.~\eqref{eq:Fassum} to \begin{equation} D(\k)=F(\k)-\sum_A\revision{\bvarepsilon^\mathrm{(true)}_{A}} f_A(\k). \end{equation} Furthermore, we assume that $F(\k)$ and $f(\k)$ are uncorrelated, which is a valid assumption for most sources of systematics since they originate from our Galaxy or due to telescope effects. Large scale surveys will reduce the current sample variance limitation on the power spectrum on scales where the systematic errors have a significant impact. As a consequence, having control of these systematics is a key requirement to provide accurate cosmological measurements. In order to investigate techniques for estimating the power spectrum in the presence of contaminants, we separate the process into three separate stages: (i) removing the contaminant signal, (ii) estimating the uncontaminated cosmological power spectrum, (iii) debiasing the resulting estimates. Two techniques are in common usage for removing the contaminant signal (i): The first is \revision{\textit{mode subtraction}} (cf. Sec.~\ref{sec:systrem}~and Sec.~\ref{sec:debias}), where contaminants are removed by fitting the amplitude of the contaminant field $f(\k)$ to the data and simply subtracted off from $F(\k)$. The second is \textit{mode deprojection} \citep{Rybicki}, which is based on assigning infinitely large covariances to contaminated modes, thus removing them from any analysis. \revision{In our nomenclature, a mode is a linear combination of Fourier modes rather than a single $\k$-mode. This is reflected in the naming of \textit{mode subtraction} and \textit{mode deprojection}. This choice of names shall distinguish the \textit{mode subtraction} technique from a third technique for removing the contaminant signal, called \textit{template subtraction}, where the observed power spectra are corrected using best-fit amplitudes derived via cross-correlations between the data and the templates. \citet*{Elsner} have shown that this method provides a biased estimate of the power and we will not consider it further in this article.} For (ii), the power spectrum P(k) is commonly estimated by the FKP estimator \citep{Feldman}, which is an approximation to the Quadratic Maximum Likelihood (QML) estimator \citep{Tegmark:1997rp}. As well as being optimal, the QML estimator has the advantage of producing unbiased power spectrum estimates. However, when applying this methodology \revision{to data with $N_\mathrm{mode}$ modes}, one has to calculate, for each bin, a $N_\mathrm{mode} \times N_\mathrm{mode}$ matrix, \revision{and then, after binning the data into $N_\mathrm{bin}$ bins,} an overall $N_\mathrm{bin} \times N_\mathrm{bin}$ normalisation matrix, which makes the application of this methodology unfeasible for future surveys with increased number of modes $N_\mathrm{mode}$. In this work, we suggest a modified FKP-style \revision{\textit{mode subtraction}} approach. We show that this technique can be made unbiased and, on a mode-by-mode basis, is mathematically identical to \textit{mode deprojection}. The FKP estimator with debiased \revision{\textit{mode subtraction}} is not optimal in that it discards more information than the full QML estimator, but we expect that, in realistic cases, this loss of information will be small. The outline of this paper is as follows: We provide an introduction to power spectrum estimation in Sec.~\ref{sec:notation}, introducing the QML and FKP estimators. We introduce the systematics removal techniques, \textit{mode deprojection} and \revision{\textit{mode subtraction}}, in sections \ref{sec:BMP} and \ref{sec:systrem}, respectively, and we show that before normalisation their resulting power spectra are the same. These are extended to multiple contaminants in Appendices \ref{sec:MPmultderiv} \&\;\ref{sec:TSmultideriv}, respectively. We introduce a new normalisation factor in Sec.~\ref{sec:debias} for a single contaminant and compare it to the normalisation of the quadratic maximum likelihood (QML) estimator of \citet{Tegmark:1997rp}. This derivation is extended to allow for a non-diagonal covariance in Appendix \ref{sec:nondiag}. We show that we can apply our methodology also to multiple contaminants in Sec.~\ref{sec:multisyst} and test the different methods on simulations in Sec.~\ref{sec:eg}. We conclude in Sec.~\ref{sec:conclusion}.
\label{sec:conclusion} We have considered methods to remove contaminants when measuring the 3D galaxy power spectrum from a given density field, focussing on \textit{mode deprojection} and \revision{\textit{mode subtraction}}. In order to understand how these are related, we have decomposed the problem into separate steps. In particular we have separated \textit{mode deprojection} from power spectrum estimation - they are often considered together - arguing that this split makes sense given the mathematical equivalence of \textit{mode deprojection} and \revision{\textit{mode subtraction}}. We argue that the QML estimation is not practical for modern surveys with large numbers of observed modes, but that we can apply \textit{mode deprojection} to the FKP-estimator, using the mathematical equivalence of \textit{mode deprojection} and \revision{\textit{mode subtraction}}, thus avoiding having to create large estimator and covariance matrices for all modes. The resulting estimate is biased, but can easily be made unbiased with a simple correction, again that can be implemented without the \revision{inversion} of large matrices. This correction is easily extended to the case of multiple contaminants and is not affected if the modes are correlated even without the effects of contaminants. The final result of our short paper is the suggestion that 3D galaxy power spectrum should be estimated using Eq.~\eqref{eq:ubPTS}, \begin{equation} \widehat{P}(k_i)=\frac{1}{N_{\Bbbk_i}}\sum_{\k_\alpha}\frac{\left\vert F(\k_\alpha)-\frac{S_P}{R_P}f(\k_\alpha)\right\vert^2}{1-\frac{1}{R_P}\frac{\vert f(\k_\alpha)\vert^2}{P(k_\alpha)}}. \end{equation} While theoretically it is sub-optimal, in practice the degradation of signal is expected to be less than ignoring window effects in the optimisation of mode averaging when using the standard FKP estimator.
16
7
1607.02417
We assess and develop techniques to remove contaminants when calculating the 3D galaxy power spectrum. We separate the process into three separate stages: (I) removing the contaminant signal, (II) estimating the uncontaminated cosmological power spectrum and (III) debiasing the resulting estimates. For (I), we show that removing the best-fitting contaminant (mode subtraction) and setting the contaminated components of the covariance to be infinite (mode deprojection) are mathematically equivalent. For (II), performing a quadratic maximum likelihood (QML) estimate after mode deprojection gives an optimal unbiased solution, although it requires the manipulation of large N_mode^2 matrices (N<SUB>mode</SUB> being the total number of modes), which is unfeasible for recent 3D galaxy surveys. Measuring a binned average of the modes for (II) as proposed by Feldman, Kaiser &amp; Peacock (FKP) is faster and simpler, but is sub-optimal and gives rise to a biased solution. We present a method to debias the resulting FKP measurements that does not require any large matrix calculations. We argue that the sub-optimality of the FKP estimator compared with the QML estimator, caused by contaminants, is less severe than that commonly ignored due to the survey window.
false
[ "mode deprojection", "recent 3D galaxy surveys", "mode subtraction", "modes", "any large matrix calculations", "infinite (mode deprojection", "FKP", "an optimal unbiased solution", "large N_mode^2 matrices", "the uncontaminated cosmological power spectrum", "contaminants", "the resulting FKP measurements", "QML", "the survey window", "rise", "a biased solution", "the resulting estimates", "II", "Peacock", "the total number" ]
12.302132
3.06941
161
1759186
[ "Bassa, C. G.", "Pleunis, Z.", "Hessels, J. W. T." ]
2017A&C....18...40B
[ "Enabling pulsar and fast transient searches using coherent dedispersion" ]
27
[ "ASTRON, the Netherlands Institute for Radio Astronomy, Postbus 2, 7990 AA, Dwingeloo, The Netherlands", "Anton Pannekoek Institute for Astronomy, University of Amsterdam, Science Park 904, 1098 XH, Amsterdam, The Netherlands", "ASTRON, the Netherlands Institute for Radio Astronomy, Postbus 2, 7990 AA, Dwingeloo, The Netherlands" ]
[ "2017ApJ...846L..19P", "2017ApJ...846L..20B", "2018A&C....25..205S", "2018ApJ...864...16M", "2018ApJ...866..149Z", "2018IAUS..337...33B", "2019A&A...623A..90S", "2019A&A...626A.104S", "2019A&ARv..27....4P", "2019ApJ...883...42N", "2019ApJ...884...14S", "2020A&A...634A...3V", "2020A&C....3000337B", "2020ApJ...896L..41C", "2020JAI.....950018F", "2021A&A...649A.120M", "2021A&A...655A..16M", "2021ApJ...911L...3P", "2021MNRAS.503.5367B", "2022A&A...657A..46B", "2022A&A...661A..87S", "2022RAA....22j5007L", "2023ApJ...958..191S", "2023ApJS..269...29N", "2023RAA....23a5023Z", "2024MNRAS.527.4397M", "2024Univ...10..158R" ]
[ "astronomy" ]
11
[ "Methods", "Data analysis-pulsars", "General", "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "1971ApJ...169..487H", "1975MComP..14...55H", "2000A&AS..147..195M", "2002AJ....124.1788R", "2003ApJ...596.1142C", "2005AJ....129.1993M", "2006Sci...314...97K", "2007Sci...318..777L", "2009IEEEP..97.1421E", "2009IEEEP..97.1497L", "2010MNRAS.406L..60N", "2010MNRAS.409..619K", "2010Natur.467.1081D", "2010PASP..122..595N", "2011A&A...530A..80S", "2011MNRAS.417.2642M", "2011PASA...28....1V", "2012MNRAS.422..379B", "2013A&A...556A...2V", "2013ITAP...61.2540E", "2013MNRAS.431.1352B", "2013PASA...30....7T", "2013Sci...340..448A", "2013Sci...341...53T", "2014A&A...570A..60C", "2014ApJ...791...67S", "2014Natur.505..520R", "2015ApJ...812...81L", "2015ApJS..218...23A", "2015Natur.528..523M", "2016A&A...585A.128K", "2016MNRAS.458.1267V", "2016MNRAS.460L..30C", "2016PASA...33...45P" ]
[ "10.1016/j.ascom.2017.01.004", "10.48550/arXiv.1607.00909" ]
1607
1607.00909_arXiv.txt
Advances in electronics, computing and networking, primarily following Moore's law, have enabled the realization of a new type of radio telescope. Instead of placing receivers at the focus of movable reflecting dishes, these new telescopes employ a large number of stationary dipole antennas to create what is called an \textit{aperture array}. The signals from these antennas are combined digitally in a correlator to create images of the sky, or in a beamformer to form beams on the sky. Three major digital aperture arrays operating at low radio frequencies are presently in operation: the LOw Frequency Array (LOFAR; \citealt{hwg+13}), the Murchison Wide Field Array (MWA; \citealt{lcm+09,tgb+13}) and the Long Wavelength Array (LWA; \citealt{ecc+09,etc+13}), while the low-frequency component of the Square Kilometre Array (SKA1-Low) is being planned \citep{bbg+15}. To maximize sensitivity while keeping the number of individual antennas, and hence cost, low, these aperture arrays operate at long wavelengths, and hence low observing frequencies (below 300\,MHz). One of the key science areas at these low observing frequencies is the study of radio pulsars. Radio emission from these highly magnetized, rotating neutron stars exhibits a very steep spectrum \citep{mkkw00,blv13}, typically peaking or turning over between 100 to 200\,MHz (e.g.\,\citealt{mgj+94}). Surveys at these low frequencies have the prospect of discovering new pulsars that are too faint to be detected in surveys performed at higher observing frequencies and can take advantage of large fields of view (e.g.\,\citealt{clh+14}). Of particular interest are radio pulsars spinning at millisecond spin periods. These millisecond pulsars provide unparalleled precision for measuring neutron star masses, performing precision tests of General Relativity, understanding binary evolution, and detecting gravitational waves (e.g.\ \citealt{ksm+06,dpr+10,rsa+14,afw+13,vlh+16}). Radio emission propagating through the ionized interstellar medium suffers from dispersion, introducing a frequency dependent time delay over the requisite large bandwidths of radio astronomical observations. As a result, pulsed signals, such as those of pulsars and fast transients (a generic term used for other sources of millisecond-duration radio pulses; e.g.\,\citealt{lbm+07}), have a specific dispersion measure ($\mathrm{DM}$) which relates directly to the column density of free electrons between the source and the observer. When searching for new pulsars and fast transients, correcting for this dispersion is typically done through \textit{incoherent dedispersion}, where the dispersive delays are removed by time shifting the time-series of individual, narrow, frequency channels by an amount appropriate to the DM of the source. Though this corrects for dispersion between channels, the dispersion within the finite bandwidth of the individual channels is not corrected for. Nonetheless, the computational efficiency of the technique has been a practical necessity compared to more accurate approaches. A priori, the DM of a new pulsar or fast transient is unknown, and the data must be dedispersed to a broad range of different DMs. Models for the Galactic electron density (e.g.\,\citealt{cl02}) can be used to estimate the maximum DM towards a given direction though -- to enable sensitivity to extra-Galactic fast transients -- most ongoing surveys search up to a maximum DM of several thousand pc\,cm$^{-3}$. Depending on the frequency and time resolution of the input data, several thousand DM trials may need to be computed \citep{cm03}\footnote{See \citealt{lsc+16b} and \url{http://www.jb.man.ac.uk/pulsar/Surveys.html} for the parameters of ongoing and historic pulsar and fast transient surveys.}. Though dedispersing that many DM trials is a computationally expensive task, recent implementations of incoherent dedispersion algorithms on graphics processing units (GPUs) are fast enough to allow real-time processing \citep{mks+11,bbbf12,slbn16}. \begin{figure} \includegraphics[angle=270,width=\columnwidth]{dmf} \caption{The effect of residual dispersion smearing in frequency channels as a function of dispersion measure (DM) and frequency. The diagonal lines denote the $3\sigma$ detection limit of an undispersed input pulse of $10\sigma$ with pulse full-width-half-maxima of $w=0.1$, 1 and 10\,ms. For frequencies below, or DMs above, these limits, denoted by the hashed areas, the pulse is no longer detectable. The channel size is set at $\Delta \nu=0.02$\,MHz which corresponds to a time resolution of $\Delta t=\Delta \nu^{-1}=50$\,$\upmu$s. This time resolution is what is typically used in current millisecond pulsar searches. Using narrower channels would adversely reduce the time resolution. The frequency bands and central frequencies of representative radio telescopes are shown with the horizontal lines.} \label{fig:dmf} \end{figure} At low frequencies and/or high DMs, incoherent dedispersion can lead to significant smearing of the pulse (in time) within a channel (see Fig.\,\ref{fig:dmf}). The effect of dispersion can be completely removed though \textit{coherent dedispersion}. This approach convolves the input signals with the inverse of the transfer function of the interstellar medium. This convolution must be performed before the signal is detected (squared) as the phase information, in addition to the amplitude, is required. Hence, the data rate and computational requirements for coherent dedispersion are typically larger than the filterbanked data used for incoherent dedispersion, as for the latter the two polarizations can be squared and frequency and/or time resolution can be reduced. Because of these higher data rates, coherent dedispersion is presently only used for observing either known pulsars or when searching for pulsars in globular clusters with known DMs. Here we present \texttt{cdmt}, for \textit{coherent dispersion measure trials}, which implements the coherent dedispersion algorithm to perform coherent dedispersion to many dispersion measure trials in parallel on GPUs. This software allows us to control the residual dispersion smearing within a channel and retain both high time and high frequency resolution when searching for pulsars and fast transients. In a semi-coherent dedispersion search, the input data can be coherently dedispersed to several coarsely separated trial DMs, each of which is then incoherently dedispersed with finer DM steps around the coherent trial DM. Though the total number of incoherent DM trials, and hence processing requirements, will increase, this approach allows us to search for millisecond pulsars at lower observing frequencies than were previously possible -- thus probing a new astrophysical parameter space. The paper is structured as follows. The coherent dedispersion algorithm, combined with channelizing the data, is described in \S\,\ref{sec:description}. Our implementation of the algorithm is outlined in \S\,\ref{sec:implementation}. In \S\,\ref{sec:application} we provide an application example of the software and we report on the performance in \S\,\ref{sec:performance}. Finally, we discuss prospects for \texttt{cdmt} in \S\,\ref{sec:discussion}.
\label{sec:discussion} Correcting for dispersion in surveys for pulsars and FRBs through incoherent dedispersion is no longer a major computational bottleneck. Recent implementations of incoherent dedispersion algorithms on GPUs are fast enough to allow real-time processing \citep{mks+11,bbbf12,slbn16}. We have presented \texttt{cdmt}, the next step in correcting for dispersion; an implementation to perform coherent dedispersion for many different DM trials. These DM trials serve as input for further incoherent dedispersion and allow semi-coherent dedispersion searches for pulsars and FRBs. The combination of coherent and incoherent dedispersion limits the dispersion smearing and allows a more flexible choice of channel size and sampling time at different observing frequencies and dispersion measures. We are using \texttt{cdmt} in an ongoing LOFAR survey of radio millisecond pulsars associated with unidentified \textit{Fermi} $\gamma$-ray sources. With this approach, we limit the dispersion smearing at these low observing frequencies over the DM range being surveyed, retaining sensitivity to short period pulsars at all DMs being sampled (see Fig.\,\ref{fig:surveys}). The success of this approach has already been demonstrated with the discovery of a new millisecond pulsar (Fig.\,\ref{fig:j1552}). Surveys for pulsars and fast transients are key science goals for SKA1-Low and SKA1-Mid \citep{hpb+15,kbk+15,mkg+15}. Though the use of semi-coherent dedispersion is currently not planned, the SKA1-Low survey for pulsars at high Galactic latitudes can benefit from this approach. As the SKA1-Low pulsar survey will likely operate at higher frequencies (200 to 350\,MHz; \citealt{kbk+15}) in comparison to LOFAR, dispersion smearing will be lower and hence fewer coherent DM trials would be required. For an example setup of 100\,MHz of bandwidth at a center frequency of 250\,MHz, a time resolution of 50\,$\upmu$s and 20\,kHz channels, coherently dedispersed DM step sizes of 10\,pc\,cm$^{-3}$ would provide a maximum dispersion smearing of 90\,$\upmu$s. Semi-coherent dedispersion searches may also be useful at higher observing frequencies. Recent FRB discoveries at observing frequencies of 1.4\,GHz have DMs in excess of 1,000\,pc\,cm$^{-3}$ \citep{tsb+13,cpk+16,pbj+16} and for these the width of the pulses is dominated by dispersion smearing within a channel (Fig.\,\ref{fig:surveys}). Coherently dedispersing to a few DM trials spaced at intervals of 10 to 100\,pc\,cm$^{-3}$ would keep the pulse broadening due to dispersion smearing below 0.1\,ms and increase the signal to noise of detected pulses, while also giving a more accurate portrayal of the intrinsic pulse duration. The \texttt{cdmt} software is continuing development, with new features being planned. Features being worked on include implementing the spectral kurtosis method \citep{ng10a,ng10b}, in order to reject radio frequency interference when computing data offsets and scales to remove its influence on the redigitization, and reading and writing different input and output formats. The code is publicly available at \url{http://github.com/cbassa/cdmt}. \begin{small}
16
7
1607.00909
We present an implementation of the coherent dedispersion algorithm capable of dedispersing high-time-resolution radio observations to many different dispersion measures (DMs). This approach allows the removal of the dispersive effects of the interstellar medium and enables searches for pulsed emission from pulsars and other millisecond-duration transients at low observing frequencies and/or high DMs where time broadening of the signal due to dispersive smearing would otherwise severely reduce the sensitivity. The implementation, called cdmt, for coherent dispersion measure trials, exploits the parallel processing capability of general-purpose graphics processing units to accelerate the computations. We describe the coherent dedispersion algorithm and detail how cdmt implements the algorithm to efficiently compute many coherent DM trials. We apply the concept of a semi-coherent dedispersion search, where coherently dedispersed trials at coarsely separated DMs are subsequently incoherently dedispersed at finer steps in DM. The software is used in an ongoing LOFAR pilot survey to test the feasibility of performing semi-coherent dedispersion searches for millisecond pulsars at 135 MHz. This pilot survey has led to the discovery of a radio millisecond pulsar-the first at these low frequencies. This is the first time that such a broad and comprehensive search in DM-space has been done using coherent dedispersion, and we argue that future low-frequency pulsar searches using this approach are both scientifically compelling and feasible. Finally, we compare the performance of cdmt with other available alternatives.
false
[ "many coherent DM trials", "semi-coherent dedispersion searches", "coherent dispersion measure trials", "coherent dedispersion", "low observing frequencies", "high DMs", "searches", "many different dispersion measures", "DM", "a semi-coherent dedispersion search", "millisecond pulsars", "DMs", "future low-frequency pulsar searches", "other available alternatives", "dispersive smearing", "finer steps", "time", "general-purpose graphics processing units", "the coherent dedispersion algorithm", "pulsed emission" ]
10.206867
3.953973
-1
12500424
[ "Tsuji, Takashi" ]
2016PASJ...68...84T
[ "Near-infrared spectroscopy of M dwarfs. IV. A preliminary survey on the carbon isotopic ratio in M dwarfs*" ]
10
[ "Institute of Astronomy, School of Science, The University of Tokyo, 2-21-1 Osawa, Mitaka-shi, Tokyo 181-0015, Japan" ]
[ "2018ASSL..451..219A", "2018ASSL..453....3D", "2019ApJ...871L...3C", "2020AJ....159...30H", "2020PASJ...72..102I", "2021A&A...656A..76Z", "2022AJ....163...72I", "2023ApJ...954..121C", "2024MNRAS.529.3171R", "2024arXiv240313760H" ]
[ "astronomy" ]
2
[ "ISM: abundances", "stars: abundances", "stars: atmospheres", "stars: late-type", "stars: low-mass", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1929PASP...41..271S", "1930PASP...42...34M", "1960stat.book..569M", "1969ApJ...157..737G", "1973ApJ...185..213D", "1974ApJ...193..621L", "1974ApJ...193..631T", "1974ApJS...28....1J", "1975ApJ...199..436T", "1975ApJ...200..675D", "1976ApJ...203..455D", "1977PASJ...29..711F", "1980ARA&A..18..399W", "1981A&A....94..175R", "1983JMoSp..98...64G", "1983JMoSp..99..431C", "1986ApJ...311..843S", "1986ApJS...62..373L", "1987MNRAS.224..237H", "1989GeCoA..53..197A", "1990ApJS...72..387S", "1992adsu.book.....W", "1993A&A...275..101E", "1993Natur.364...25O", "1993Natur.365..806A", "1994ARA&A..32..191W", "1996A&A...309..760P", "1996A&A...310..933O", "1997MNRAS.289L..11A", "2000SPIE.4008.1056K", "2002A&A...396..967P", "2002ApJ...573L.137A", "2002ApJ...575..264T", "2005A&A...441..181C", "2006ApJS..165..618A", "2006MNRAS.368.1087B", "2008A&A...477..865S", "2008A&A...489.1271T", "2009ARA&A..47..481A", "2010JQSRT.111.2139R", "2013ARA&A..51..457N", "2014PASJ...66...98T", "2015PASJ...67...26T", "2016PASJ...68...13T" ]
[ "10.1093/pasj/psw076", "10.48550/arXiv.1607.07004" ]
1607
1607.07004_arXiv.txt
Besides the elemental abundances, isotopic ratios in stars and interstellar matter (ISM) provide important clues on stellar evolution and Galactic chemical evolution. However, isotopic effects are prominent only in molecular spectra but not in atomic spectra. For this reason, isotopic analysis has mostly been done on molecular spectra observed in cool celestial objects such as cool stars and the ISM. Because of this limitation, our knowledge on the isotopic abundances in the Universe has been rather poor compared to that on the elemental abundances. On the other hand, spectroscopic analyses of the isotopic ratios are somewhat simpler, compared to those of the elemental abundances, in that the isotopic ratios are relatively insensitive to the physical condition of the environments where spectral lines are formed. For this reason, isotopic analyses have extensively been done on the spectra of cool stars since the middle of the 20-th century even when analyses on the elemental abundances were more difficult. For example, initial attempt to determine the $^{12}$C/$^{13}$C ratio in cool giant stars has been done by \citet{Gre69}, who showed that the $^{12}$C/$^{13}$C ratio in red giant stars is decreased compared to the solar system value of 89.9 \citep{And89}. Then, extensive analyses on the $^{12}$C/$^{13}$C ratio in G and K giant stars have been done with the use of photoelectric scans of the CN red system (\cite{Day73}; \cite{Tom74}; \cite{Tom75}; \cite{Dea75}). The results confirmed that the products of the CNO cycle are dredged-up in the red giant phase, but the resulting $^{12}$C/$^{13}$C ratios appeared to be too small compared with the theoretical predictions (e.g., \cite{Dea76}). Studies on the $^{12}$C/$^{13}$C ratio were extended to red supergiants \citep{Lam74} and to AGB stars with the use of the high resolution FTS spectra (\cite{Smi86}, \yearcite{Smi90}; \cite{Tsu08}). The results on oxygen-rich giants, supergiants, and AGB stars revealed that these stars contribute to lower the $^{12}$C/$^{13}$C ratio in the ISM by their mass-loss. Historically, it was on carbon stars that studies on carbon isotopes in stellar spectra have initiated: It was nearly a century ago when $^{13}$C bearing molecules were identified by the C$_2$ Swan bands, first as $^{12}$C$^{13}$C \citep{San29} and next as $^{13}$C$^{13}$C \citep{Men30} in carbon stars (also referred to as R and N type stars). Initial works suggested that the $^{12}$C/$^{13}$C ratios in most carbon stars are in the range of 2 -- 20, except for a few exceptional case (\cite{Mck60} and references cited therein). This result suggested that the $^{12}$C/$^{13}$C ratios in carbon stars may be related to the equilibrium value of the $^{12}$C/$^{13}$C ratio in the CN cycle. However, since the $^{12}$C$^{12}$C Swan bands in carbon stars are very strong and heavily saturated while the $^{12}$C$^{13}$C Swan bands less saturated, the intensity ratio of $^{12}$C$^{12}$C/$^{12}$C$^{13}$C should apparently be small leading to a low $^{12}$C/$^{13}$C ratio. For this reason, the very low $^{12}$C/$^{13}$C ratios found in carbon stars can be due to imperfect correction for the saturation effect, and a possibility that $^{12}$C/$^{13}$C $\approx 100$ was suggested for some cool carbon stars (e.g., \cite{Fuj77}). Since then, many works on the $^{12}$C/$^{13}$C ratio have been carried out for a large sample of carbon stars: \citet{Lam86} determined the $^{12}$C/$^{13}$C ratios in dozens of N type stars through the analysis of the high resolution FTS spectra of the CN red system, CO first and second overtones, and showed that the $^{12}$C/$^{13}$C ratios in cool carbon stars extend to as large as the solar system value. Determinations of the $^{12}$C/$^{13}$C ratios were extended through the analysis of the CN red system to 62 N type stars by \citet{Ohn96} and to 44 carbon stars by \citet{Abi97}. The results on carbon-rich AGB stars revealed that the $^{12}$C/$^{13}$C ratios in these stars are generally larger than in oxygen-rich AGB stars and these stars contribute to increase the $^{12}$C/$^{13}$C ratio in the ISM by their mass-loss. In contrast, isotopic analyses on cool dwarfs are quite limited, even though extensive analyses on elemental abundances in late-type dwarfs have been done (e.g., \cite{Edv93}). It is true that molecular lines are not so strong in F and G dwarfs. However, the major reasons for this contrast between high and low luminosity stars may be that the high resolution spectroscopy needed to isolate faint isotopic features could be applied more easily to bright high luminous stars on one hand and also some isotopic features appeared to be enhanced in high luminous cool stars due to the evolutionary effects in the stars themselves. In M dwarfs, determinations of the elemental abundances have also been difficult, but we have shown recently that the carbon and oxygen abundances in M dwarfs could be determined rather accurately by the use of the near infrared spectra of CO and H$_2$O, respectively (\cite{Tsu14}, \cite{Tsu15}, \cite{Tsu16}, hereafter Papers I, II, and III, respectively). This has been made possible by several reasons: First, the well known difficulty of the continuum in cool stars has been overcome in doing spectroscopic analysis by referring to the pseudo-continuum, which can be evaluated accurately with the use of the recent molecular line-list including many weak lines of H$_2$O (e.g., \cite{Bar06}; \cite{Rot10}). Then, spectroscopic analysis can be done essentially the way as referring to the true-continuum, so long as the pseudo-continua can be defined consistently in the observed and predicted spectra. Second, given that carbon and oxygen abundances are determined from stable molecules such as CO and H$_2$O, respectively, the problem of photospheric model should not be so serious. This is because CO consumes most of carbon and H$_2$O most of oxygen left after CO formation, and CO abundance is almost identical with that of C and H$_2$O abundance with that of O$-$C in M dwarfs. For this reason, CO and H$_2$O abundances are insensitive to the uncertainties of the photospheric structures \footnote{ It is to be noted that this advantage does not apply to CO in general, but applies to CO in M dwarfs only. In warmer F or G dwarfs, CO abundance is highly sensitive to temperature and cannot be used for accurate abundance analysis (see footnote 9 in Paper I).}, and carbon and oxygen abundances could be determined rather well despite the use of model photospheres which anyhow cannot be very accurate (as for further details, see e.g., Paper II). Given that the elemental abundances of at least some elements were well determined in M dwarfs, a further important possibility in M dwarfs is that the isotopic ratios can also be determined from the molecular spectra, at the same time as the elemental abundances in the same object. In the near infrared spectra of M dwarfs, not only $^{12}$C$^{16}$O but also $^{13}$C$^{16}$O bands are observed, and the $^{12}$C/$^{13}$C ratios can be discussed. A problem is that the spectral resolution of our spectra is about $\lambda/\Delta\,\lambda \approx 20000 $ or the velocity resolution is about 16\,km\,s$^{-1}$. This resolution barely made it possible to measure equivalent widths and hence to determine the elemental abundances of carbon and oxygen (Papers I, II, and III). However, this resolution is not high enough to measure line profiles accurately and hence to analyze faint isotopic features. For this reason, our analysis on the carbon isotopic ratios is necessarily preliminary. Even with this limitation, however, we hope that we can discuss the elemental and isotopic abundances of carbon simultaneously in a large sample of M dwarfs for the first time. In this paper, we first summarize the known data on our program stars from Papers I, II, and III in section 2. We then investigate how to estimate the $^{12}$C/$^{13}$C ratios from the spectra of $^{13}$CO in section 3. The resulting $^{12}$C/$^{13}$C ratios, with assessments of the accuracy, are given in section 4. The resulting carbon isotopic ratios in M dwarfs are discussed in comparisons with those of the ISM in section 5.
We have tried to estimate the carbon isotopic ratio $^{12}$C/$^{13}$C in M dwarfs based on the spectra of medium resolution ($\lambda/\Delta\,\lambda \approx 20000$). Although this resolution is certainly not sufficient to resolve faint $^{13}$C$^{16}$O lines from the stronger lines of $^{12}$C$^{16}$O and H$_2$O, we find clear evidence for the $^{13}$C$^{16}$O feature in the spectra of M dwarfs for the first time (see Fig.\,1 and Figs.\,2a - l). Then, determination of the $^{12}$C/$^{13}$C ratios in M dwarfs is quite feasible especially if higher spectral resolution can be employed, and we hope that the isotopic analysis on low luminosity stars such as M dwarfs will be done more intensively as in high luminosity cool stars. With our medium resolution, we determine a possible best value of $^{12}$C/$^{13}$C in each M dwarf by comparing observed and predicted spectra, not by a direct inspection but by an application of the chi-square method. Although the direct comparisons (see Figs.\,2a - l, 3a-f) show difficulties inherent to the spectra of insufficient resolution, the chi-square analysis works reasonably well (see Fig.\,3). As a result, we determine $^{12}$C/$^{13}$C ratios for 31 M dwarfs and a lower limit of 200 for 17 M dwarfs. From our preliminary survey, we may suggest that the $^{12}$C/$^{13}$C ratios in M dwarfs are larger than about 40 and not likely to be so small as in some evolved high luminosity stars. Although the lower limit of the $^{12}$C/$^{13}$C ratio is found to be 200 in many M dwarfs, this result is uncertain by a factor of two or so. Unfortunately, it is beyond the capability of our medium resolution spectra to analyze faint features due to very low abundance of $^{13}$C. We hope that our result can be improved by the use of higher spectral resolutions in the near future. The mean $^{12}$C/$^{13}$C ratio in M dwarfs turns out to be larger than that of the ISM (section\,5). While the $^{12}$C/$^{13}$C ratio in the ISM provides the present $^{12}$C/$^{13}$C ratio, that in M dwarfs reflects the carbon isotopic ratio in the past ISM. So far, the solar system, conserving the $^{12}$C/$^{13}$C ratio of the ISM 4.5 Gyr ago, was used as a reference in investigating the time variation of the $^{12}$C/$^{13}$C ratio in the ISM, but M dwarfs can in principle be used for the same purpose with numerous objects in a larger time-space domain. In fact, M dwarfs offer an unique possibility to determine the $^{12}$C/$^{13}$C ratios in unevolved late-type dwarfs and should have important role to trace the evolution of carbon isotopes (and possibly other isotopes) in the Galaxy. For this reason, determination of isotopic abundances in M dwarfs should be the next major challenge when high resolution infrared spectroscopy will be ripe enough to be able to observe many faint objects not well observed so far. \bigskip I thank T. Nakajima for recovering the spectra observed at slit positions A and B before co-adding and for helpful discussion on the evaluation of the S/N ratio. I also thank him and Y. Takeda for sharing the spectra of M dwarfs observed at Subaru IRCS. I thank the anonymous referee for careful reading of the text and for critical comments especially on the accuracy of the analysis. This research has made use of the SIMBAD operated at CDS, Strasbourg, France. Computations are carried out on common use data analysis computing system at the Astronomical Data Center, ADC, of the National Astronomical Observatory of Japan.
16
7
1607.07004
Carbon isotopic ratios are estimated in 48 M dwarfs based on the medium resolution near infrared spectra (λ/Δ λ ≈ 20000) of the <SUP>13</SUP>CO (3,1) band. We find clear evidence for the presence of a <SUP>13</SUP>CO feature for the first time in the spectra of M dwarfs. Spectral resolution of our observed data, however, is not high enough to analyze the <SUP>13</SUP>CO feature directly. Instead, we compare the observed spectrum with synthetic spectra assuming <SUP>12</SUP>C/<SUP>13</SUP>C = 10, 25, 50, 100, and 200 for each of 48 M dwarfs and estimate the best possible <SUP>12</SUP>C/<SUP>13</SUP>C ratio by chi-square analysis. The resulting <SUP>12</SUP>C/<SUP>13</SUP>C ratios in M dwarfs distribute from 39 to a lower limit of 200. The mean value of 31 M dwarfs for which <SUP>12</SUP>C/<SUP>13</SUP>C ratios are determined (i.e., excluding those with the lower limit only) is (<SUP>12</SUP>C/<SUP>13</SUP>C)<SUB>dM</SUB> = 87 ± 21 (p.e.), and that of 48 M dwarfs including those with the lower limit of 200 is (<SUP>12</SUP>C/<SUP>13</SUP>C)<SUB>dM</SUB> &gt; 127 ± 41 (p.e.). These results are somewhat larger than the <SUP>12</SUP>C/<SUP>13</SUP>C ratio of the present interstellar matter (ISM) determined from the molecular lines observed in the millimeter and optical wavelength regions. Since the amount of <SUP>13</SUP>C in the ISM has increased with time due to mass loss from evolved stars, the <SUP>12</SUP>C/<SUP>13</SUP>C ratios in M dwarfs, reflecting those of the past ISM, should be larger than those of the present ISM. In M dwarfs, log <SUP>13</SUP>C/<SUP>12</SUP>C plotted against log A<SUB>C</SUB> shows a large scatter without clear dependence on the metallicity. This result shows a marked contrast to log <SUP>16</SUP>O/<SUP>12</SUP>C (= log A<SUB>O</SUB>/A<SUB>C</SUB>) plotted against log A<SUB>C</SUB>, which shows a rather tight correlation with a larger value at the lower metallicity. Such a contrast can be a natural consequence of <SUP>16</SUP>O and <SUP>12</SUP>C being primary products in stellar nuclear synthesis while <SUP>13</SUP>C is a secondary product, at least partly.
false
[ "M dwarfs", "SUP>12</SUP", "M", "p.e.", "dwarfs", "SUP>13</SUP", "ISM", "SUP>12</SUP>C/<SUP>13</SUP>C ratio", "stellar nuclear synthesis", "Carbon isotopic ratios", "primary products", "SUP>12</SUP>C/<SUP>13</SUP>C", ".", "evolved stars", "31 M dwarfs", "48 M dwarfs", "C</SUB", "mass loss", "Δ", "synthetic spectra" ]
7.372174
12.32756
-1
12472918
[ "Bernal, C. G.", "Fraija, N." ]
2016MNRAS.462.3646B
[ "A central compact object in Kes 79: the hypercritical regime and neutrino expectation" ]
3
[ "Instituto de Matemática, Estatística e Física, Universidade Federal do Rio Grande, Av. Itália km 8 Bairro Carreiros, Rio Grande, 96203-900 RS, Brazil", "Instituto de Astronomía, Universidad Nacional Autónoma de México, Circuito Exterior, C.U., A. Postal 70-264, 04510 México D.F., Mexico" ]
[ "2017JPhCS.932a2006D", "2018PASP..130l4201F", "2018PhRvD..98h3012F" ]
[ "astronomy" ]
3
[ "equation of state", "magnetic fields", "MHD", "neutrinos", "stars: neutron", "supernovae: individual: Kes 79", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1968AuJPh..21..739K", "1972PhRvD...6..941D", "1978PhRvD..17.2369W", "1982PhR....90...73B", "1989ApJ...341..867C", "1989ApJ...346..847C", "1989neas.book.....B", "1993ApJ...411..823W", "1995APh.....3..267L", "1995ApJ...440L..77M", "1996ApJS..102..411I", "1996PhRvL..77.3082A", "1998ApJ...504..761C", "1998PhRvL..81.1774A", "1999A&A...345..847G", "1999JCoPh.154..284P", "2000ApJS..126..501T", "2000ApJS..131..273F", "2001ApJ...548..800C", "2002PhRvD..66a3001C", "2003ApJ...584..414S", "2003RvMP...75..345G", "2004IAUS..218..239P", "2004JHEP...04..078A", "2004mnpa.book.....M", "2005ApJ...619..839C", "2005ApJ...627..390G", "2008PhR...460....1G", "2010PNAS..107.7147K", "2010PhRvD..81i2004W", "2010RMxAA..46..309B", "2011MNRAS.414.2567H", "2011PPN....42..577G", "2011PhRvL.107x1801A", "2012ARNPS..62...81S", "2012ApJ...748..148S", "2012MNRAS.425.2487V", "2013ApJ...770..106B", "2013PhRvC..88b5501A", "2014ApJ...783...44F", "2014ApJ...787..140F", "2014ApJ...790...94B", "2014MNRAS.437.2187F", "2014MNRAS.441.1209F", "2014MNRAS.442..239F", "2015APh....70...54F", "2015APh....71....1F", "2015MNRAS.450.2784F", "2015MNRAS.451..455F", "2015PASA...32...18P", "2015PhRvC..91c5806R", "2016ApJ...826...31F", "2016JHEAp...9...25F", "2016MNRAS.456.3813T" ]
[ "10.1093/mnras/stw1874", "10.48550/arXiv.1607.05652" ]
1607
1607.05652.txt
% %, typical sizes of 0.3 - 3 km Central Compact Objects (CCOs) are point-like sources located at central regions of some supernova remnants, which have been observed only in the X-rays \citep{2004IAUS..218..239P}. The X-ray spectra have been characterized to have thermal components with blackbody temperatures of 0.2 - 0.5 keV and luminosities $L_X\sim 10^{33} - 10^{34}\, {\rm erg\, s^{-1}}$. Several of such sources and their supernova remnants (SNRs) are well-known, including RCW103, Cassiopeia A (Cas A), Pup A, and Kes 79 \citep{2010PNAS..107.7147K}. Currently, it is widely accepted that CCOs are neutron stars born in supernova explosions with an unusually estimated low magnetic field. A possible explanation for this atypical behavior is the so-called hidden magnetic field scenario. It suggests that the strong magnetic field has been buried due to a strong accretion during a core-collapse supernova (see e.g., \cite{1995ApJ...440L..77M}, \cite{1999A26A...345..847G}, \cite{Bernal2013} and references therein). When the core-collapse supernova event takes place, the shock is still pushing its way through the outer layers of the progenitor, and if it encounters a density discontinuity, a reverse shock may be generated. Depending on its strength and how far it was generated, this reverse shock can induce a strong accretion onto the newborn neutron star on a timescales of hours. Hypercritical accretion results when the accretion rate is higher than Eddington accretion rate ($\dot{M}>\dot{M}_{Edd}$). In this scenario, photons are trapped within the accretion flow and the energy liberated by the accretion is lost through neutrino emission close to the neutron star surface. After the reverse shock hits the neutron star surface and rebounds, a third shock develops and starts moving outward against the in-falling matter. Once this accretion shock stabilizes it will separate the in-falling matter from an extended envelope in quasi-hydrostatic equilibrium. Based on such scenario of late accretion onto newborn neutron stars inside supernovae, \cite{1989ApJ...341..867C} developed an analytical model for the hypercritical regime. In this model, the neutrino cooling plays an important role in the formation of a quasi-hydrostatic envelope around the compact remnant.\\ % In \cite{1989ApJ...346..847C}, the author highlighted some conditions involved in the formation of neutron stars inside supernovae, including the hyperaccretion of material due reverse shock and the formation of a convective envelope around the compact remnant. Following such ideas, \cite{1995ApJ...440L..77M} considered the submergence of the magnetic field on the stellar crust and the subsequent ohmic diffusion of the submerged magnetic flux through the outermost nonmagnetic layers of the crust of a newborn neutron star. With these requirements, \cite{1999A26A...345..847G} presented simple 1D ideal magnetohydrodynamical (MHD) simulations of the effect of this post-supernova hypercritical accretion on the newborn neutron star magnetic field to show that such submergence/re-emergence could occur. These simulations displayed a rapid submergence of the magnetic field into the neutron star. Based on MHD simulations with high refinement, \cite{Bernal2010,Bernal2013} confirmed this result. After hyperaccretion stops, the magnetic field could diffuse back to the surface, resulting in a delayed switch-on of a pulsar. It is worth noting that depending on the amount of accreted material, the submergence could be so deep that the neutron star may appear and remain unmagnetized for centuries \citep{1999A26A...345..847G}. Recently, some authors revisiting this scenario studied the magnetic field evolution in the CCOs context \citep{2011MNRAS.414.2567H, 2012MNRAS.425.2487V}. In addition, \cite{Shabaltas2012} and \cite{Popov2015} have suggested that such hidden magnetic field scenario may be applicable for various CCOs, including the compact remnant in Kes 79. \\ %. Due this similarity, it is proposed that this CCO is a neutron star with a hidden magnetic field due the hyperaccretion. Located in the Galactic plane at 33 degrees northeast of the Galactic center, Kes 79 (also known as G33.6+0.1) is the source 79 in the radio catalog of Kesteven \citep{Kesteven1968}. It is a moderately large SNR with a point-like source at its center, widely believed to be a neutron star created in the SN explosion. Such CCO is called CXO J1852.6+0040 \citep{Seward2003}. The authors showed that the luminosity, in the X-ray band (0.3-8 keV), is $\sim 7\times10^{33}$ erg s$^{-1}$, which corresponds to four times the X-ray luminosity reported in the CCO of Cas A. The blackbody spectrum peaking at X-rays and the lack of a Pulsar Wind Nebula (PWN) indicate that the thermal emission could be originated from a small region, perhaps on the surface of the neutron star. As \cite{Popov2015} pointed out, the hypothesis of their suppressed magnetosphere is mainly based on the analysis of their thermal emission: pulse profiles of the X-ray light curves and a high pulse fraction, which requires magnetar-scale fields in the crust.\\ % \cite{2005ApJ...627..390G} reported the discovery of 105 ms X-ray pulsations from the CCO in Kes 79, with an upper limit on its spin-down rate of $\dot P<7 \times 10^{-14}$ s s$^{-1}$. Assuming a magnetic dipole model they estimate that the surface magnetic field strength is $B<3 \times 10^{12}$ G. Also, if a blackbody model of temperature $T_{BB}=(0.44 \pm 0.03$) keV is used for the X-ray spectrum characterization, they estimated a radius for the source of $R_{BB} \sim 0.9$ km. More recently, \cite{2014ApJ...790...94B} modeled such X-ray pulsations in Kes 79 in the context of thermal surface radiation from a rotating neutron star and more accurate results for the surface magnetic field were estimated. Taking into account the reasons mentioned above and the similarities between the CCO in Kes 79 and the neutron star in Cas A, in the present work we adopt, for simplicity, standard parameters for the neutron star in Kes 79: $M \simeq 1.44$ $M_{\odot}$, $R \simeq10$ km, and an average pre-hyperaccretion magnetic field of $B \simeq 10^{12}$ G. Such parameters are adequate for progenitor models of pre-supernova in the range 20--40 $M_{\odot}$ as it seems to be the case of Kes 79 (see \cite{Chevalier2005, Shabaltas2012} and references therein.) Although we have chosen the Kes 79 CCO as a particular case, the method presented here can be adapted for other CCOs with similar features. \\ % In the hidden magnetic field scenario, the neutrino production and cooling on the neutron star surface play an important role in the formation of a quasi-hydrostatic envelope \citep{Bernal2013, 2014MNRAS.442..239F, 2015MNRAS.451..455F}. The properties of these neutrinos get modified when they propagate in this thermal and magnetized medium. For instance, depending on the flavor composition, neutrinos would feel a different effective potential. These changes in the flavor mixing are due to the electron neutrino ($\nu_e$) interaction with electrons via both neutral and charged currents (CC), whereas muon ($\nu_\mu$) and tau ($\nu_\tau)$ neutrinos interact only via the neutral current (NC). The resonant conversion of neutrino from one flavor to another caused by the medium effect is well-known as the Mikheyev-Smirnov-Wolfenstein effect \citep{wol78}.\\ % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Using the FLASH code to simulate the dynamics close to the stellar surface (including the reverse shock, the neutrino cooling processes and the magnetic field dynamics), some authors presented numerical studies of hypercritical accretion of matter onto the neutron star surface in the center of the SNRs 1987A and Cas A \citep{2014MNRAS.442..239F,2015MNRAS.451..455F}. The authors estimated the number of neutrinos that must have been seen from the hypercritical accretion episode on the Super-Kamiokande (SK) and Hyper-Kamiokande (HK) neutrino experiments. In this work, we study the dynamics of the envelope and analyze the submergence of the magnetic field into the stellar surface for the compact remnant in Kes 79. We estimate the neutrino flux and the flavor ratio expected on Earth. The paper is arranged as follows. In section \ref{sec-Physics} we describe the physics included in the hypercritical model. In section \ref{sec-FLASH} we show the numerical results from MHD simulations. In section \ref{sec-Neutrinos} we study the thermal neutrino dynamics and we analyze the neutrino oscillations. Finally, a brief conclusions are drawn in section \ref{sec-Results}. % % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\label{sec-Results} % In the present work we have studied the dynamics of hyperaccretion (with neutrino emission) onto a newly born neutron star as is the case of Kes 79 scenario. To do this, we have performed numerical 2D-MHD simulations using the AMR FLASH method. The code allows to capture the rich morphology and the main characteristics of the fundamental physical processes in detail. We included a custom routine in the code that take into account several neutrino cooling processes, which are active depending of strict criteria of temperature and density of the model in the computational domain. With this, we found that for the estimated hyperaccretion rate for the Kes 79 scenario, the bulk of magnetic field is submerged efficiently into the new crust formed by the accreted matter during this regime. Additionally, we found that the code reproduce the radial profile of the main thermodynamical parameters of the hypercritical regime, including the accretion shock radius, which it is in good agreement with the value analytically estimated. It is noteworthy that a (completely relaxed) steady state could not be reached, within the parameters of the simulation, due to various factors: periodicity of the lateral boundaries, effects of buoyancy and magnetic stresses near the stellar surface, among others. Also, although most of the magnetic field is confined/submerged in the stellar crust, residual weak magnetic field is present within the quasi-hydrostatic envelope but that does not influence the dynamics of the system. Once the hypercritical phase ends, the story for the magnetic field continues. \cite{1995ApJ...440L..77M} found that after hyperaccretion stopped, the bulk of the submerged magnetic field could diffuse back to the surface by ohmic processes and depending on the amount of accreted matter, the submergence could be so deep that the neutron star may appear and remain unmagnetized for centuries. It is possible that, in addition to the ohmic diffusion of the bulk of the magnetic field, some portion of the confined magnetic field may be pushed into the neutron star by the hyperaccretion, and it may be crystallized in the mantle. The presence of a crystal lattice of atomic nuclei in the crust is mandatory for modeling of the subsequent radio-pulsar glitches (see \cite{1999A26A...345..847G} and references therein). Presence of solid crust enables excitation of toroidal modes of oscillations. The toroidal modes in a completely fluid star have all zero frequency, but the presence of a solid crust gives them nonzero frequencies in the range of kHz. % On the other hand, requiring the value of distance $d_z=7.$1 kpc \citep{1998ApJ...504..761C}, neutrino Luminosity $L_{\bar{\nu}_e}=(8.4\pm 0.4)\times10^{48}\:\mathrm{erg\, s^{-1}}$ (this work, Fig. \ref{fig6}), effective volume of HK, $V\simeq 0.56\times 10^{12}\,{\rm cm^3}$ \citep{2014arXiv1412.4673H} and the average neutrino energy $<E_{\bar{\nu}_e}>\simeq 7\, {\rm MeV}$, from eq. (\ref{num_Neu}) we obtain that the number of events that could have been expected from the hypercritical phase on Hyper-Kamiokande is 733$\pm$364. In addition, we compute the number of initial neutrino burst expected during the neutron star formation. Taking into account a temperature $T \approx 4\pm1$ MeV \citep{Giunti-Kim-2007}, a duration of the neutrino pulse in the hypercritical phase $t\simeq10^{3}$ s, the average neutrino energy $<E_{\bar{\nu}_e}>\simeq 13.5\pm 3.2\, {\rm MeV}$ \citep{Giunti-Kim-2007} and a total fluence equivalent of Kes 79 is $\Phi\approx(1.25\pm 0.52) \times 10^{11}\, \bar{\nu}_e \,{\rm cm^{-2}}$ \citep{2004mnpa.book.....M, 1989neas.book.....B}, then the total number of neutrinos emitted from Kes 79 would be $N_{tot}=6\,\Phi\, 4\pi\, d_z^2\approx(8.97\pm 3.58) \times 10^{57}$. Similarly, we can compute the total radiated luminosity corresponding to the binding energy of the neutron star $L_\nu\approx \frac{N_{tot}}{t}\times <E_{\bar{\nu}_e}>\approx (2.16 \pm0.87) \times 10^{52}$ erg/s. Regarding the effective volume $V\simeq 0.56\times 10^{12}\,{\rm cm^3}$ \citep{2014arXiv1412.4673H}, from eq. (\ref{num_Neu}) we get $N_{ev}\simeq 1129\pm 475$ events expected during the neutron star formation in a neutrino detector as HK experiment. Comparing the number of neutrinos expected during the neutron star formation and the hyperaccretion phase, we obtain 1.5 events.\\ % Neutrinos generated at the hypercritical accretion phase will oscillate in their ways due to electron density in zones I, II, III and IV, and after in vacuum into Earth. In zone I, the thermal plasma is endowed with a magnetic field $\sim$ (1 - 6)$\times 10^{12}$ G and thermalized $\sim$ (1 - 8) MeV. Taking into account the neutrino effective potential given in zone I, we plot the neutrino effective potential as a function of magnetic field, angle, temperature and chemical potential. From fig. \ref{fig4} can be observed the positivity of the effective potential ($V_{eff}>$ 0), therefore neutrinos can oscillate resonantly. Taking into account the two and three- neutrino mixing parameters we analyze the resonance condition, as shown in Figure \ref{fig5}. Recently, \citet{2014MNRAS.442..239F} showed that neutrinos can oscillate resonantly due to the density profiles of the collapsing material surrounding the progenitor. In zone II, III and IV, the author showed that neutrinos can oscillate resonantly and computed the survival and conversion probabilities for the active-active ($\nu_{e,\mu,\tau} \leftrightarrow \nu_{e,\mu,\tau}$) neutrino oscillations in each region. Taking into account the oscillation probabilities in each region and in the vacuum (on its path to Earth), we calculate the flavor ratio expected on Earth for neutrino energies of $E_{\nu}=1$ MeV, 1 MeV, 3 MeV, 5 MeV and 7 MeV, as shown in table \ref{Table}. In this table we can see a small deviation from the standard ratio flavor 1:1:1. In this calculation we take into account that for neutrino cooling processes (electron-positron annihilation, inverse beta decay, nucleonic bremsstrahlung and plasmons), only inverse beta decay is the one producing electron neutrino. It is worth noting that our calculations of resonant oscillations were performed for neutrinos instead of anti-neutrinos, due to the positivity of the neutrino effective potential.\\ As a final remark, we can say that the present numerical studies of these phenomena are necessary and very important and its results can give us a glimpse of the complex dynamics around the newly born neutron stars, moments after the core-collapse supernova explosion. Although we use the Kes 79 as a particular case of the hidden magnetic field scenario, this method can be applied to the various CCOs scenarios. Nevertheless, the eventual growing of the bulk magnetic field post-hyperaccretion, the possible magnetic field amplification by turbulent dynamo and the magnetic field crystallization processes inside neutron star require more detailed studies that are outside the scope of this paper. % % \begin{table*} \begin{center} \caption[]{The neutrino flavor ratio in each region of the hypercritical phase for $E_{\nu}=$ 1, 3, 5 and 7 MeV.}\label{Table} \begin{minipage}{126mm} \begin{tabular}{lccccccccc} \hline $E_{\nu}$ & On the NS surface & Accretion material & Free fall zone& Outer layers & On Earth \\\hline \hline {\small 1} & {\small 1.2:0.9:0.9} & {\small 1.186:0.907: 0.907} & {\small 1.152:0.924:0.924} & {\small 1.120:0.940:0.940} & {\small 1.036:0.988:0.976} \\\hline {\small 3} & {\small 1.2:0.9:0.9} & {\small 1.161:0.919:0.919}& {\small 1.130:0.935:0.935} & {\small 1.132:0.934: 0.934} & {\small 1.039:0.987:0.974}\\\hline {\small 5} & {\small 1.2:0.9:0.9} & {\small 1.159:0.921:0.921} & {\small 1.127:0.937:0.937} & {\small 1.123:0.938:0.938} & {\small 1.037:0.987:0.975} \\\hline {\small 7} & {\small 1.2:0.9:0.9} & {\small 1.172: 0.914:0.914} & {\small 1.142:0.929:0.929} & {\small 1.135:0.933:0.933} & {\small 1.041:0.987:0.973} \\\hline \end{tabular} \end{minipage} \end{center} \end{table*} %
16
7
1607.05652
We present magnetohydrodynamical simulations of a strong accretion on to magnetized proto-neutron stars for the Kesteven 79 (Kes 79) scenario. The supernova remnant Kes 79, observed with the Chandra ACIS-I instrument during approximately 8.3 h, is located in the constellation Aquila at a distance of 7.1 kpc in the galactic plane. It is a galactic and a very young object with an estimate age of 6 kyr. The Chandra image has revealed, for the first time, a point-like source at the centre of the remnant. The Kes 79 compact remnant belongs to a special class of objects, the so-called central compact objects (CCOs), which exhibits no evidence for a surrounding pulsar wind nebula. In this work, we show that the submergence of the magnetic field during the hypercritical phase can explain such behaviour for Kes 79 and others CCOs. The simulations of such regime were carried out with the adaptive-mesh-refinement code FLASH in two spatial dimensions, including radiative loss by neutrinos and an adequate equation of state for such regime. From the simulations, we estimate that the number of thermal neutrinos expected on the Hyper-Kamiokande Experiment is 733 ± 364. In addition, we compute the flavour ratio on Earth for a progenitor model.
false
[ "such regime", "such behaviour", "radiative loss", "CCOs", "Kes", "objects", "thermal neutrinos", "a surrounding pulsar wind nebula", "neutrinos", "state", "Aquila", "Chandra", "the galactic plane", "instrument", "The Kes 79 compact remnant", "Kes 79 and others CCOs", "magnetized proto-neutron stars", "two spatial dimensions", "the adaptive-mesh-refinement code FLASH", "an adequate equation" ]
4.538416
4.746485
60
12623583
[ "De Felice, Antonio", "Mukohyama, Shinji" ]
2017PhRvL.118i1104D
[ "Graviton Mass Might Reduce Tension between Early and Late Time Cosmological Data" ]
34
[ "Center for Gravitational Physics, Yukawa Institute for Theoretical Physics, Kyoto University, 606-8502 Kyoto, Japan", "Center for Gravitational Physics, Yukawa Institute for Theoretical Physics, Kyoto University, 606-8502 Kyoto, Japan and Kavli Institute for the Physics and Mathematics of the Universe (WPI), UTIAS, The University of Tokyo, 277-8583 Chiba, Japan" ]
[ "2013arXiv1309.2146M", "2017EPJC...77..725W", "2017JCAP...04..039M", "2017JPhCS.883a2009D", "2017PhRvD..95l3540D", "2017PhRvD..96b3506H", "2017PhRvD..96b4032D", "2017PhRvD..96d4029M", "2017PhRvD..96j4036D", "2017arXiv170904141W", "2018PhRvD..98b4010B", "2018PhRvD..98d3527N", "2018PhRvD..98f4037D", "2018PhRvD..98j4031D", "2019JCAP...01..017A", "2019PhRvD..99d4055D", "2019PhRvD..99h3509B", "2019PhRvD..99j3526K", "2020EPJC...80..708A", "2020PhRvD.101l3521B", "2021A&A...653A.148H", "2021JCAP...04..015D", "2021JCAP...04..018D", "2021JCAP...12..011D", "2021MNRAS.505.5427N", "2021PhRvD.104h4013F", "2021PhRvD.104j4057D", "2021arXiv210503849K", "2022PhRvD.105j4013D", "2022PhRvD.105l3026M", "2022PhRvD.106h4050D", "2023IJTP...62..184W", "2023PTEP.2023g2E01A", "2023PhRvD.107f4070M" ]
[ "astronomy", "physics" ]
2
[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Theory" ]
[ "1974ITAC...19..716A", "2004MNRAS.353.1201P", "2006PhRvD..74l3507T", "2008Natur.451..541G", "2009JCAP...10..004S", "2010CAMCS...5...65G", "2010PhRvD..82d4020D", "2011JCAP...11..030E", "2011MNRAS.415.2876B", "2011PhRvL.106w1101D", "2012ApJ...751L..30H", "2012JCAP...03..006E", "2012MNRAS.420.2102S", "2012MNRAS.423.3430B", "2012MNRAS.424.2339T", "2012MNRAS.425..405B", "2012MNRAS.426.2581G", "2012PhRvD..85l3503K", "2012PhRvL.109q1101D", "2013A&A...557A..54D", "2013PASP..125..306F", "2013PhRvL.111p1301M", "2015MNRAS.449..848H", "2016A&A...594A..13P", "2016JCAP...04..028D", "2016MNRAS.460.4188G", "2016MNRAS.461.3781C", "2016PASJ...68...38O", "2016PhLB..752..302D", "2016PhRvL.116f1102A", "2016PhRvL.116v1101A" ]
[ "10.1103/PhysRevLett.118.091104", "10.48550/arXiv.1607.03368" ]
1607
1607.03368_arXiv.txt
16
7
1607.03368
The standard cold dark matter model with a cosmological constant (Λ -CDM) predicts a growth of structures which tends to be higher than the values of redshift space distortion (RSD) measurements if the cosmological parameters are fixed by the cosmic microwave background data. In this Letter, we point out that this discrepancy can be resolved or understood if we assume that the graviton has a small but nonzero mass. In the context of the minimal theory of massive gravity (MTMG), due to infrared Lorentz violations measurable only at present cosmological scales, the graviton acquires a mass without being haunted by unwanted extra degrees of freedom. While the so-called self-accelerating branch of cosmological solutions in the MTMG has the same phenomenology for the background as well as the scalar- and vector-type linear perturbations as the Λ CDM in general relativity (GR), it is possible to choose another branch so that the background is the same as that in GR, but the evolution of matter perturbations gets modified by the graviton mass. In studying the fit of such modified dynamics to the above-mentioned RSD measurements, we find that the Λ CDM model is less probable than the MTMG by 2 orders of magnitude. With the help of the cross-correlation between the integrated Sachs-Wolfe effect and the large-scale structure, the data also pin down the graviton mass squared around μ<SUP>2</SUP>≈-(3 ×10<SUP>-33</SUP> eV )<SUP>2</SUP>, which is consistent with the latest bound |μ<SUP>2</SUP>|&lt;(1.2 ×10<SUP>-22</SUP> eV )<SUP>2</SUP> set by the recent LIGO observation.
false
[ "SUP>2</SUP>|&lt;(1.2", "present cosmological scales", "graviton", "cosmological solutions", "<SUP>2</SUP>≈-(3 ×10", "matter perturbations", "unwanted extra degrees", "general relativity", "freedom", "μ", "GR", "MTMG", "infrared Lorentz violations", "the cosmic microwave background data", "the graviton mass", "structures", "such modified dynamics", "magnitude", "RSD", "the cosmological parameters" ]
9.916084
0.759942
-1
901416
[ "Henson, Monique A.", "Barnes, David J.", "Kay, Scott T.", "McCarthy, Ian G.", "Schaye, Joop" ]
2017MNRAS.465.3361H
[ "The impact of baryons on massive galaxy clusters: halo structure and cluster mass estimates" ]
87
[ "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" ]
[ "2016JCAP...11..038K", "2016arXiv161204247P", "2017ApJ...849...54L", "2017MNRAS.465..213B", "2017MNRAS.465.2936M", "2017MNRAS.471.1088B", "2018ApJS..235...20H", "2018JCAP...02..005H", "2018MNRAS.474.3173P", "2018MNRAS.477.4957B", "2018MNRAS.478.2618F", "2018MNRAS.479..890L", "2018MNRAS.479.5385H", "2018PASJ...70S..28M", "2019ApJ...884...33G", "2019MNRAS.482.1352M", "2019MNRAS.482.3308A", "2019MNRAS.483.3459R", "2019MNRAS.484...60M", "2019MNRAS.484.1526A", "2019MNRAS.484.1946G", "2019MNRAS.488.3340F", "2019MNRAS.489.2439H", "2019MNRAS.489.3565J", "2019MNRAS.490.4889H", "2020A&A...634A.113A", "2020A&A...644A.126T", "2020ApJ...901....5B", "2020MNRAS.491.1575C", "2020MNRAS.491.1622P", "2020MNRAS.492.2285D", "2020MNRAS.492.4780O", "2020MNRAS.493.3274L", "2020MNRAS.495..784S", "2020MNRAS.495.1737H", "2020MNRAS.496.2743G", "2020MNRAS.496.4468S", "2020MNRAS.497.1332B", "2020MNRAS.497.4684H", "2020MNRAS.499.3445Y", "2020OJAp....3E..13C", "2020PhRvD.102b3509A", "2021A&A...652A.155S", "2021ApJ...908...91H", "2021ApJ...910...28W", "2021ApJ...910...32C", "2021MNRAS.500.5056R", "2021MNRAS.502.5115G", "2021MNRAS.503.5310R", "2021MNRAS.504.5383D", "2021MNRAS.505..593D", "2021MNRAS.505.3907L", "2021MNRAS.506.2533B", "2021MNRAS.508.6092A", "2022A&A...666A..34S", "2022ApJ...936..161W", "2022AstL...48....1K", "2022JCAP...10..034C", "2022MNRAS.509.1127S", "2022MNRAS.509.3441A", "2022MNRAS.513.2178H", "2022MNRAS.513.4754F", "2022MNRAS.514..313B", "2022MNRAS.514.2876X", "2022MNRAS.515.3383D", "2022MNRAS.515.4471W", "2022MNRAS.515.6023D", "2022MNRAS.517.5303L", "2023A&A...671A..57S", "2023A&A...674A..80L", "2023MNRAS.518.4238G", "2023MNRAS.520.4000F", "2023MNRAS.520.5845T", "2023MNRAS.520.6223P", "2023MNRAS.524.2262A", "2023PhDT.........8A", "2023arXiv230409142Q", "2023arXiv230613187S", "2024A&A...681A..67E", "2024A&A...681A..87R", "2024A&A...682A..31B", "2024A&A...682A.157L", "2024MNRAS.528.3797W", "2024MNRAS.tmp.1430B", "2024PhRvL.132c1001L", "2024arXiv240408539K", "2024arXiv240508557K" ]
[ "astronomy" ]
14
[ "gravitational lensing: weak", "galaxies: clusters: general", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
[ "1965TrAlm...5...87E", "1985ApJ...292..371D", "1997ApJ...490..493N", "1998ApJ...495...80B", "1998ApJ...499L...5M", "2000ApJ...529L..69J", "2000ApJ...534...34W", "2001ApJ...555..240B", "2001ApJ...556L..91S", "2001ApJ...557..533F", "2001MNRAS.328..726S", "2001PhR...340..291B", "2003MNRAS.338...14P", "2004MNRAS.349.1039N", "2004MNRAS.354...10M", "2005ApJ...618....1H", "2005ApJ...627..647B", "2005Natur.435..629S", "2005astro.ph.10346T", "2006AJ....132.2685M", "2006ApJ...640..691V", "2007ApJ...655...98N", "2007MNRAS.376..215B", "2007MNRAS.378...55M", "2007MNRAS.381.1450N", "2008MNRAS.383.1210S", "2008MNRAS.387..536G", "2008MNRAS.387..998M", "2008MNRAS.387.1431D", "2008MNRAS.390L..64D", "2008MNRAS.391.1940M", "2009A&A...498..361P", "2009ApJ...693.1142S", "2009MNRAS.393...99W", "2009MNRAS.394L..11S", "2009MNRAS.398...53B", "2009MNRAS.398.2177L", "2010ApJ...711.1033Z", "2010ApJ...721..875O", "2010MNRAS.402.1536S", "2010MNRAS.404.1137B", "2010MNRAS.405.2161D", "2010SPIE.7741E..1SN", "2011ARA&A..49..409A", "2011ApJ...740...25B", "2011ApJ...740..102K", "2011ApJS..197...30Z", "2011MNRAS.411..584M", "2011MNRAS.412.1965M", "2011MNRAS.414.1851O", "2011MNRAS.415L..69J", "2012ApJ...745L...3L", "2012ApJ...756..128F", "2012ApJS..199...25P", "2012MNRAS.420.3303B", "2012MNRAS.421..621B", "2012MNRAS.421.1073B", "2012MNRAS.422.1999K", "2012MNRAS.422.3081M", "2012MNRAS.427.1322L", "2012NJPh...14e5018R", "2013ATel.5032....1R", "2013ApJ...767..116M", "2013ApJ...769L..35O", "2013ApJ...779..112N", "2013AstRv...8a..40R", "2013MNRAS.429.3316B", "2013MNRAS.436..697B", "2014A&A...564A.129I", "2014A&A...571A..29P", "2014ApJ...797...34M", "2014MNRAS.439....2V", "2014MNRAS.439...48A", "2014MNRAS.439.2485C", "2014MNRAS.440.2290M", "2014MNRAS.441.1270L", "2014MNRAS.441.3359D", "2014MNRAS.442.2641V", "2014MNRAS.443.1973V", "2014SPIE.9153E..1PB", "2015AAS...22544304S", "2015ApJ...799..214B", "2015MNRAS.449..685H", "2015MNRAS.451.1247S", "2015MNRAS.451.1460K", "2015MNRAS.452.1171W", "2016A&A...592A...4L", "2016A&A...594A..24P", "2016ApJ...827..112B", "2016MNRAS.456L..74S", "2016MNRAS.457.4063S", "2016MNRAS.457.4340K", "2016MNRAS.459.2973S", "2017MNRAS.465..213B", "2017MNRAS.465.2584B", "2017MNRAS.465.2936M", "2017MNRAS.470..166H", "2019ApJ...873..111I" ]
[ "10.1093/mnras/stw2899", "10.48550/arXiv.1607.08550" ]
1607
1607.08550_arXiv.txt
Galaxy clusters are a sensitive probe of the late time evolution of the Universe, providing crucial insights into the nature of both dark matter and dark energy. Cluster based cosmological tests require well constrained masses for large samples of clusters. There is a long standing debate about the bias in X-ray cluster masses \citep[see][]{Mazzotta2004,Rasia2012,Applegate2014,Smith2016}, which arises due to the assumption that clusters are in hydrostatic equilibrium. Since clusters are often unrelaxed systems, this is frequently not a valid assumption. Instead, many authors are moving towards using masses derived from weak lensing (WL) observations of clusters~\citep{Okabe2010,Mahdavi2013,Hoekstra2015,Kettula2015} or at the very least, calibrating X-ray masses using weak lensing measurements~\citep[e.g.][]{Lieu2015}. The power of cluster counting has been highlighted in Sunyaev-Zel'dovich surveys performed by the \textit{Planck} satellite~\citep{Planck2015} and the South Pole Telescope~\citep{Bocquet2015}, however more accurate cluster mass measurements are needed for cluster cosmology to be competitive with other techniques~\citep{Allen2011,Planck2015}. High quality observational data is forthcoming with the ongoing and upcoming Dark Energy Survey~\citep{DES2005}, SPT-3G~\citep{Benson2014}, Large Synoptic Sky Survey~\citep{Ivezic2008} and ACTpol~\citep{Niemack2010}, but we also need simulations to provide robust theoretical predictions for comparison, as well mock data for testing observational techniques. Galaxy clusters have been extensively studied in dark matter only (DMO) simulations. It is well established in those simulations that cold dark matter haloes are triaxial, prolate structures. The sphericity of dark matter haloes decreases with increasing mass, so that galaxy clusters typically have sphericities of $(c/a){\simeq}0.4{-}0.6$~\citep{Maccio2008,Muonoz-Cuartas2011,Bryan2013}. Since both concentration and spin have also been shown to decrease weakly with mass~\citep{Bett2007,Duffy2008,Klypin2011,Muonoz-Cuartas2011,Ludlow2012,Klypin2016}, high mass clusters typically have low concentrations and exhibit little rotational support. DMO simulations have also been instrumental in testing observational methods for measuring cluster masses. Weak gravitational lensing provides a promising method for measuring the masses of galaxy clusters, since it does not require any assumptions about the dynamical state of the cluster. DMO simulations have shown that weak lensing masses are typically biased low by ${\sim}5\%$, with this bias decreasing with increasing mass~\citep{Oguri2011,Becker2011,Bahe2012}. Understanding this bias is crucial for cluster cosmology, since it requires large samples of clusters with accurately determined masses. Cosmological hydrodynamic simulations have shown that including baryons can have a significant effect upon the mass distribution of groups and low-mass clusters~\citep[e.g.][]{Bryan2013,Velliscig2014,Cusworth2014,Schaller2015}. The inclusion of baryonic effects in cosmological simulations leads to the depletion of high mass clusters~\citep{Cusworth2014}, and the clusters that do form are more spherical and have higher concentrations than their dark matter only counterparts~\citep{Duffy2010,Bryan2013}. The baryon fraction and hence the total mass within clusters is sensitive to galaxy formation processes~\citep{Stanek2009,McCarthy2011,Martizzi2012,Velliscig2014,LeBrun2014}. Thus, the impact of baryons on the shape and density profile of clusters depends on galaxy formation efficiency~\citep{Bryan2013,Duffy2010}. The impact of baryons on the mass distribution of low-mass clusters is not just limited to the central regions of clusters; feedback from Active Galactic Nuclei (AGN) can alter low-mass cluster profiles out to $R_{200}$~\citep{Velliscig2014}\footnote{$M_\Delta$ is defined as the mass contained within a sphere of radius $R_\Delta$, at which the enclosed average density is $\Delta$ times the critical density of the Universe}. It is still unclear what effect baryons will have on high-mass clusters. If baryons have a significant impact on the mass distribution of massive galaxy clusters, this may have implications for mass estimation techniques such as cluster weak lensing, which have been tested on dark matter only simulations~\citep{Oguri2011,Becker2011,Bahe2012}. The lack of hydrodynamic simulations of massive galaxy clusters is a natural consequence of the large computational cost of such simulations. Furthermore, accounting for baryonic effects is not a trivial task, requiring calibrated models for star formation, feedback from supernovae and AGN, and radiative cooling. Cosmological zoom simulations, in which the region of interest in simulated at a higher resolution than the surrounding region, offer a solution to this problem. This approach has been used on cluster scales \citep[e.g.][]{Martizzi2014a,Hahn2015}, however it has only been applied to small numbers of clusters to date. This places limitations on the conclusions of such work, since the dynamic range in mass needed to investigate mass dependent properties is lacking and it is difficult to determine whether any results are significant or an artefact of the small sample size. To obtain a sample sufficiently large to investigate the properties of massive galaxy clusters, we combine the $400\,h^{-1}\mathrm{Mpc}$ BAryons and HAloes of MAssive Systems (BAHAMAS) simulation~\citep{McCarthy2016} with the hydrodynamic zoom simulations that were developed as part of the MAssive ClusterS and Intercluster Structures (MACSIS) project~\citep{Barnes2016}. The paper is organised as follows. The simulations used and the methods used to identify haloes and classify relaxed structures are described in Section 2. In Section 3 the methods used to measure the spins, shapes and density profiles of clusters are outlined and results are presented. This is followed by the results from a weak lensing analysis of the cluster sample in Section 4. In Section 5 we discuss hydrostatic bias in this cluster sample and the method used to calculate the X-ray hydrostatic masses. Finally, we summarise our results in Section 6.
In this study we have used the MACSIS and BAHAMAS simulations presented in~\cite{Barnes2016} and~\cite{McCarthy2016} to create a combined sample of more than 3,500 clusters with $M_{200}{\geq}5{\times}10^{13}\,h^{-1}\mathrm{M}_\odot$, simulated with realistic baryonic physics. These simulations have been shown to reproduce the observed scalings of gas mass, integrated Sunyaev-Z'eldovich signal and X-ray luminosity with mass, as well as the observed hot gas radial profiles of clusters at $z=0$~\citep{McCarthy2016,Barnes2016}. We focus our study on three key areas: the properties of high mass clusters and the impact of baryons upon them, the influence of baryonic effects upon weak lensing mass estimates, and the mass dependence of the hydrostatic bias in high mass clusters. Since the MACSIS simulations consist of matched HYDRO and DMO zoom simulations, we are able to directly compare clusters simulated with and without baryonic effects. We also investigated the redshift dependence of our results. Our main results are as follows: \begin{itemize} \item The distributions of spins in the HYDRO and DMO simulations are consistent with each other and are well fitted by a lognormal distribution. The dark matter component has a slightly larger spin in the HYDRO simulations than in the DMO simulations, which is associated with a transfer in angular momentum from the baryonic component to the dark matter. Spin declines weakly with mass at all redshifts considered here (Fig.~\ref{fig:spin-shape-mass}). The slope is consistent between the HYDRO and DMO simulations and is unchanged for a relaxed subsample. The mean spin of relaxed haloes is 15\% smaller than for the entire cluster sample. \item Clusters in the HYDRO simulations are more spherical, with larger values of $s$ and $e$ on average. The sphericity-mass relation is steeper in the DMO simulations, but the elongation-mass relation is consistent between the DMO and HYDRO simulations. A larger mass range is required to constrain the effect of baryons on this slope. Selecting only relaxed haloes does not affect the slope of either the sphericity-mass or elongation-mass relation. \item By matching MACSIS clusters in the DMO and HYDRO simulations we demonstrated that clusters in the HYDRO simulations are more concentrated in the central regions (Fig.~\ref{fig:macsis-dmo-hydro-profdiff}). This is partly due to the condensation of baryons in the cluster centre and also a consequence of the contraction of the dark matter halo in the presence of baryons. The dark matter density profiles of MACSIS clusters at $z=0$ in the HYDRO simulations are more dense at $r{<}0.6R_{200}$ than in the DMO simulations. At $0.6{<}r/R_{200}{<}3$ the dark matter density profile is less dense in the HYDRO simulations than in the DMO simulations. At the high-mass end of the concentration mass relation ($M_{200}>10^{15}h^{-1}\mathrm{M}_\odot$) this manifests itself as an increase in concentrations in the HYDRO simulations (Fig.~\ref{fig:c200-nfw-mass}). Since the concentration-mass relation is flatter in the HYDRO simulations, clusters with masses $M_{200}{\approx}10^{15}h^{-1}\mathrm{M}_\odot$ have larger concentrations than clusters in the DMO simulations. \item The density profiles of clusters considered here are better fit by the Einasto profile than the NFW profile. This leads to a smaller bias in the masses calculated from fits to the spherically averaged density profile, with the NFW model underpredicting masses by 22\% on average and the Einasto model underpredicting masses by 8\% (Fig.~\ref{fig:einasto-nfw-3d}). Whilst cluster shear profiles are better fit by the Einasto model rather than the NFW model, this only results in a $2{-}3$\% improvement in weak lensing cluster mass estimates in the HYDRO simulations, despite the cost of adding an additional degree of freedom (Fig~\ref{fig:einasto-nfw-2d}). \item Baryons have a more significant effect on the shear profiles of clusters than on their convergence profiles. The shear profiles of HYDRO clusters are up to 15\% larger than clusters in the DMO simulations at $r{<}0.5h^{-1}\mathrm{M}_\odot$, as a consequence of the sensitivity of the shear to the central cluster region (Fig.~\ref{fig:wl-shearprof-baryons}). Despite this, the weak lensing mass bias is consistent between the DMO and HYDRO simulations, with both data sets showing that weak lensing underestimates cluster masses by ${\approx}10\%$ for clusters with $M_{200}{\leq}10^{15}h^{-1}\mathrm{M}_\odot$ and that this bias tends to zero at higher masses (Fig.~\ref{fig:wl-mbias-baryons}). \item The hydrostatic mass bias, $1-b=M_\mathrm{500,X-ray}/M_\mathrm{500,WL}$ declines from $0.8$ to $0.6$ for clusters with masses increasing from $M_{500}=10^{14}h^{-1}\mathrm{M}_\odot$ to $M_{500}=10^{15}h^{-1}\mathrm{M}_\odot$ when using X-ray hydrostatic masses calculated from spectroscopic temperature and density profiles (Fig.~\ref{fig:hydrostatic-bias-mass}). The X-ray and weak lensing masses are measured independently. We find no evidence for any redshift dependence. The mass dependence is mostly due to the spectroscopic temperature measurements (Fig~\ref{fig:hydrostatic-bias-temps-density}) that are biased low by the presence of cooler, X-ray emitting gas in the cluster outskirts. Using the true temperature and density profiles gives $b{\approx}0.04{-}0.14$ at the masses considered here, with no clear mass dependence. \item The hydrostatic bias is smaller for more spherical clusters that have a small centre of mass offset and fewer substructures, which motivates the morphological selection of clusters in X-ray surveys (Fig.~\ref{fig:hsebias-correlations}). \end{itemize} In conclusion, we find baryons have only a minor effect on the spins, shapes and weak lensing mass estimates of massive galaxy clusters. Baryons have a small effect on cluster density profiles at small radii, which is also apparent in their weak lensing shear profiles. When using spectroscopic temperatures and densities, the hydrostatic bias decreases as a function of mass, leading to a bias of ${\approx}40\%$ for high mass clusters. Further work is needed to clarify the cause of this large bias and to reconcile it with observational results.
16
7
1607.08550
We use the BAHAMAS (BAryons and HAloes of MAssive Systems) and MACSIS (MAssive ClusterS and Intercluster Structures) hydrodynamic simulations to quantify the impact of baryons on the mass distribution and dynamics of massive galaxy clusters, as well as the bias in X-ray and weak lensing mass estimates. These simulations use the subgrid physics models calibrated in the BAHAMAS project, which include feedback from both supernovae and active galactic nuclei. They form a cluster population covering almost two orders of magnitude in mass, with more than 3500 clusters with masses greater than 10<SUP>14</SUP> M<SUB>⊙</SUB> at z = 0. We start by characterizing the clusters in terms of their spin, shape and density profile, before considering the bias in both weak lensing and hydrostatic mass estimates. Whilst including baryonic effects leads to more spherical, centrally concentrated clusters, the median weak lensing mass bias is unaffected by the presence of baryons. In both the dark matter only and hydrodynamic simulations, the weak lensing measurements underestimate cluster masses by ≈10 per cent for clusters with M<SUB>200</SUB> ≤ 10<SUP>15</SUP> M<SUB>⊙</SUB> and this bias tends to zero at higher masses. We also consider the hydrostatic bias when using both the true density and temperature profiles, and those derived from X-ray spectroscopy. When using spectroscopic temperatures and densities, the hydrostatic bias decreases as a function of mass, leading to a bias of ≈40 per cent for clusters with M<SUB>500</SUB> ≥ 10<SUP>15</SUP> M<SUB>⊙</SUB>. This is due to the presence of cooler gas in the cluster outskirts. Using mass weighted temperatures and the true density profile reduces this bias to 5-15 per cent.
false
[ "cluster masses", "massive galaxy clusters", "clusters", "weak lensing mass estimates", "higher masses", "masses", "X-ray spectroscopy", "the median weak lensing mass bias", "mass", "=", "hydrodynamic simulations", "densities", "the cluster outskirts", "MAssive ClusterS", "a cluster population", "the hydrostatic bias", "MAssive Systems", "baryons", "z", "Intercluster Structures" ]
12.405185
4.337144
-1
12406354
[ "Drummond, B.", "Tremblin, P.", "Baraffe, I.", "Amundsen, D. S.", "Mayne, N. J.", "Venot, O.", "Goyal, J." ]
2016A&A...594A..69D
[ "The effects of consistent chemical kinetics calculations on the pressure-temperature profiles and emission spectra of hot Jupiters" ]
126
[ "Astrophysics Group, University of Exeter, EX4 4QL, Exeter, UK", "Astrophysics Group, University of Exeter, EX4 4QL, Exeter, UK; Maison de la Simulation, CEA-CNRS-INRIA-UPS-UVSQ, USR 3441 Centre d'étude de Saclay, 91191, Gif-Sur-Yvette, France", "Astrophysics Group, University of Exeter, EX4 4QL, Exeter, UK; Univ. Lyon, ENS de Lyon, Univ. Lyon1, CNRS, CRAL, UMR 5574, 69007, Lyon, France", "Astrophysics Group, University of Exeter, EX4 4QL, Exeter, UK; Department of Applied Physics and Applied Mathematics, Columbia University, New York, NY, 10025, USA; NASA Goddard Institute for Space Studies, New York, NY, 10025, USA", "Astrophysics Group, University of Exeter, EX4 4QL, Exeter, UK", "Intituut voor Sterrenkunde, Katholieke Universiteit Leuven, Celestijnenlaan 200D, 3001, Leuven, Belgium", "Astrophysics Group, University of Exeter, EX4 4QL, Exeter, UK" ]
[ "2016A&A...595A..36A", "2017A&A...598A..97A", "2017A&A...605A..95Y", "2017ApJ...841...30T", "2017ApJ...850...46T", "2017ApJS..228...20T", "2017MNRAS.472.2334G", "2017Natur.548...58E", "2018A&A...612A.105D", "2018A&A...618A..63B", "2018AJ....155...29W", "2018AJ....156..283E", "2018AJ....156..298A", "2018ApJ...853..138B", "2018ApJ...854....8B", "2018ApJ...855L..31D", "2018ApJ...863..183K", "2018ApJ...869...28D", "2018MNRAS.474.1705N", "2018MNRAS.474.5158G", "2018Natur.557..526N", "2018exha.book.....P", "2019A&A...624A..58V", "2019AJ....157..170M", "2019ARA&A..57..617M", "2019ApJ...871...56M", "2019ApJ...871..158F", "2019ApJ...873...32M", "2019ApJ...876..144T", "2019ApJ...880...14S", "2019ApJ...883..194M", "2019MNRAS.482.4503G", "2019MNRAS.486..783G", "2019MNRAS.486.1123D", "2019MNRAS.487.2082L", "2019MNRAS.487.2242H", "2019MNRAS.488.1332L", "2019MNRAS.488.2222M", "2019MolPh.117.3922C", "2019arXiv190303997G", "2020A&A...633A...2D", "2020A&A...636A..68D", "2020A&A...637A..38P", "2020A&A...639A..48S", "2020A&A...643A..23T", "2020AJ....159....5S", "2020AJ....160...51A", "2020ApJ...890..176V", "2020ApJ...895...77B", "2020ApJ...901..110P", "2020JQSRT.25507228T", "2020MNRAS.496.1638M", "2020MNRAS.497.5155W", "2020MNRAS.498.4680G", "2020RAA....20...99Z", "2020SSRv..216...58C", "2021A&A...646A..21C", "2021A&A...648A.127B", "2021A&A...652A..57K", "2021AJ....161..131D", "2021AJ....161..278W", "2021AJ....162..108F", "2021AJ....162..138R", "2021AJ....162..168S", "2021Ap&SS.366...83Y", "2021ApJ...923..242G", "2021ApJ...923..264T", "2021ApJ...923..269K", "2021MNRAS.502.5643L", "2021MNRAS.505.4515R", "2021MNRAS.505.5603B", "2021MNRAS.506.4500C", "2021PSJ.....2..106F", "2021exbi.book...18M", "2022A&A...658A..41F", "2022A&A...658A..42P", "2022A&A...667A..13Z", "2022A&A...667A..15K", "2022AJ....163...42Y", "2022AJ....163..190F", "2022ApJ...925L...3F", "2022ApJ...932...20D", "2022ApJ...934...31F", "2022ExA....53..279M", "2022MNRAS.512.4877B", "2022MNRAS.515.3037N", "2022MNRAS.517.1407C", "2022MNRAS.517.2383B", "2022NatAs...6..471M", "2023A&A...675A..25H", "2023ARep...67.1329Z", "2023ApJ...954...22L", "2023ApJ...956L..32G", "2023ApJ...957..104M", "2023ApJ...958...68S", "2023MNRAS.518.2472R", "2023MNRAS.519.3129Z", "2023MNRAS.520.3867T", "2023MNRAS.521.2607T", "2023MNRAS.522..648C", "2023MNRAS.523.5681N", "2023MNRAS.524..817T", "2023MNRAS.524..835R", "2023MNRAS.524.3396N", "2023MNRAS.525.1375W", "2023Natur.614..649J", "2023Natur.614..653A", "2023Natur.614..664A", "2023Natur.614..670F", "2023Natur.617..483T", "2023RASTI...2...45C", "2024A&A...681A...3G", "2024A&A...683A.194F", "2024A&A...685A.125F", "2024A&A...686A..24B", "2024AJ....167...45R", "2024ApJ...961..210P", "2024ApJ...963...73M", "2024ApJ...965..115L", "2024ApJ...966...39S", "2024ApJ...967..132A", "2024MNRAS.529.1776Z", "2024MNRAS.529.2686L", "2024MNRAS.531.1056R", "2024MNRAS.tmp.1409C", "2024arXiv240200756M" ]
[ "astronomy" ]
9
[ "planets and satellites: atmospheres", "planets and satellites: composition", "Astrophysics - Earth and Planetary Astrophysics" ]
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[ "10.1051/0004-6361/201628799", "10.48550/arXiv.1607.04062" ]
1607
1607.04062_arXiv.txt
Despite the discovery of ever smaller and more Earth-like rocky exoplanets \citep[e.g.][]{Berta-Thompson2015}, hot Jupiters remain one of the most important classes of exoplanet due to the availability of follow-up characterisation observations. This allows for important comparisons between the various observables (transmission spectra \citep[e.g.][]{Sing2016}, emission spectra \citep[e.g.][]{Knutson2008,Beaulieu2010,Diamond-Lowe2014,Evans2015} and phase curves \citep[e.g.][]{Zellem2014}; see \citet{HengShowman2015} for a review) with the many atmosphere models in use throughout the community to understand the physical and chemical processes occuring in these atmospheres. There exists a hierarchy of atmosphere models which have so far been applied to the study of exoplanets with each class of model having its own practical use. The three--dimensional (3D) general circulation models (GCMs) are required to gain insights into the atmospheric dynamics \citep[e.g.][]{Showman2009,Burrows2010,Heng2011,RauscherMenou2012,Dobbs-Dixon2013,Mayne2014,Polichtchouk2014} whilst the one--dimensional (1D) radiative--convective models solve for single column pressure--temperature ($P$--$T$) profiles \citep[e.g.][]{Iro2005,Barman2005,Fortney2010,SpiegelBurrows2010}. A further class of atmosphere model, the chemical kinetics models, determine precise chemical compositions \citep[e.g.][]{Moses2011,Venot2012,Zahnle2014}. Early work using these chemical kinetics models focussed on the photochemical production and destruction processes of hydrocarbon hazes \citep{Liang2004,Zahnle2009a} and the production of atomic hydrogen through water photolysis \citep{Liang2003}. Later modelling efforts provided detailed chemical compositions of these atmospheres and considered the impact on the observable properties \citep{Moses2011,Venot2012,Madhusudhan2011}. The non-equilibrium processes of vertical mixing and photochemistry have been shown to have notable consequences on the simulated spectra of some hot exoplanet atmospheres \citep[see][for a review]{Moses2014}. Vertical mixing can result in transport-induced quenching which increases or decreases the abundances of chemical species compared with their chemical equilibrium profiles. In addition, at low pressures high-energy photons can dissociate molecules into highly reactive daughter products. These non-equilibrium processes have the potential to influence the simulated spectrum of the atmosphere by changing the amount of absorption or emission from these molecules. For example, \citet{Moses2011} found that increases in the abundances of methane and ammonia due to transport-induced quenching have important consequences for the emission spectrum of HD 189733b, decreasing the calculated eclipse depth over several wavelength regions. One common feature amongst all of the chemical kinetics studies published so far is that none include a consistent approach to the model atmosphere. The background thermal structure of the atmosphere is treated as a fixed model input. Usually the supplied $P$--$T$ profiles have been derived from 3D GCM results \citep[e.g.][]{Moses2011,Venot2012} or 1D radiative-convective models \citep[e.g.][]{Moses2013,Agundez2014b}. The equilibrium and non-equilibrium chemical abundances are then calculated using this fixed $P$--$T$ profile leading to an inconsistency between the chemical abundances and the thermal structure. As the chemistry is driven away from chemical equilibrium, the opacity in the model atmosphere changes which should impact on the $P$--$T$ profile. \citet{Agundez2014b} have attempted to take into account this effect by including one additional iteration in their radiative-convective model using the initial non-equilibrium chemical abundances they calculate for the hot Neptune GJ 436b. They found corrections to the temperature structure of $<$ 100 K. In addition, \citet{Hubeny2007} performed consistent non-equilibrium chemistry models of brown dwarf atmospheres using a timescale argument to quench the abundances of N$_2$/NH$_3$ and CO/CH$_4$, not using chemical kinetics, and found corrections to the $P$--$T$ profile of 50--100 K. In this study we use a 1D atmosphere code which solves for both hydrostatic and energy balance and includes a sophisticated chemistry scheme to solve for the non-equilibrium chemical abundances, including vertical mixing and photochemistry, consistently with the $P$--$T$ profile. In \cref{section:model_description} we describe our model setup. In \cref{section:results} we present our results and, finally, in \cref{section:conclusion} we summarise and conclude.
\label{section:conclusion} We have presented results based on a fully-consistent chemical kinetics model applied to the atmospheres of HD~189733b and HD~209458b. Our simulations show that in cases of strong disequilibrium chemistry transport-induced quenching of absorbing species can induce changes in the $P$--$T$ profile of up to 100 K. These temperature shifts can, in turn, have impacts on both the chemical abundances themselves and on the corresponding emission spectra. The chemical abundances are affected via two related processes: firstly, temperature changes at high pressures, where the chemistry remains in chemical equilibrium, induce new chemical equilibrium abundances, and secondly, the change in temperature shifts the quenching point which alters the quenched abundances at lower pressures. For instance, in our model of HD~189733b we would conclude that CH$_4$ is more abundant than CO for the $K_{zz}$ = 10$^{11}$ cm$^{2}$s$^{-1}$ case when not performing the calculation consistently. Instead, in the consistent approach, we find that CO is the most abundant carbon-bearing species. For the model of HD~209458b (with a temperature inversion) we find that the abundances of CH$_4$ and NH$_3$ are a factor of $\sim$5 and $\sim$3.5 lower in the consistent model compared with the non-consistent model, due to an increase in the temperature by more than 100 K in the deep atmosphere. For both HD~189733b and HD~209458b models we find that consistent calculations of non-equilibrium chemistry reduce the overall impact of chemical disequilibrium on the emission spectrum. Our results show that in conventional chemical kinetics models, where the $P$--$T$ profile is held fixed, the dominant mechanism for non-equilibrium chemistry to affect the emission spectrum is by changing the pressure level, and temperature, of the photosphere. The strong dependence of the emission flux on temperature will result in a very different simulated emission spectrum. However, in our consistent model the $P$--$T$ profile adapts to the new non-equilibrium chemical composition and to retain energy balance in the model atmosphere. The consequent temperature changes mitigate the effect of changing the location of the photosphere by either heating up or cooling down at the location of the new photosphere to preserve energy balance. Based on these results, we urge caution when assessing the impact of non-equilibrium chemistry (transport-induced quenching and photochemistry) on the emission spectrum. Not including consistency between the chemical abundances and the temperature structure can lead to overestimates of the impact of non-equilibrium chemistry. This work has only consisdered 1D (vertical) effects of non-equilibrium chemistry. Horizontal advection is expected to be very important in the atmospheres of tidally-locked exoplanets which possess very strong zonal wind velocities \citep{Showman2009,Heng2011,RauscherMenou2012,Mayne2014}. In addition, these atmospheres can possess very large day-night temperature contrasts leading to large contrasts in horizontal chemical equilibrium abundances\citep{Burrows2010,Kataria2016}. Disequilibrium chemistry has already been suggested as a possible explanation to explain the discrepencies between the observed and model emission phase curves \citep{Zellem2014}. However, this work has shown that when performed consistently, transport-induced quenching has a smaller impact on the emission spectrum than previous studies suggest. It would therefore be very interesting and timely to study the process considered here including horizontal advection, by coupling a chemical kinetics scheme consistently to a 3D GCM.
16
7
1607.04062
In this work we investigate the impact of calculating non-equilibrium chemical abundances consistently with the temperature structure for the atmospheres of highly-irradiated, close-in gas giant exoplanets. Chemical kinetics models have been widely used in the literature to investigate the chemical compositions of hot Jupiter atmospheres which are expected to be driven away from chemical equilibrium via processes such as vertical mixing and photochemistry. All of these models have so far used pressure-temperature (P-T) profiles as fixed model input. This results in a decoupling of the chemistry from the radiative and thermal properties of the atmosphere, despite the fact that in nature they are intricately linked. We use a one-dimensional radiative-convective equilibrium model, ATMO, which includes a sophisticated chemistry scheme to calculate P-T profiles which are fully consistent with non-equilibrium chemical abundances, including vertical mixing and photochemistry. Our primary conclusion is that, in cases of strong chemical disequilibrium, consistent calculations can lead to differences in the P-T profile of up to 100 K compared to the P-T profile derived assuming chemical equilibrium. This temperature change can, in turn, have important consequences for the chemical abundances themselves as well as for the simulated emission spectra. In particular, we find that performing the chemical kinetics calculation consistently can reduce the overall impact of non-equilibrium chemistry on the observable emission spectrum of hot Jupiters. Simulated observations derived from non-consistent models could thus yield the wrong interpretation. We show that this behaviour is due to the non-consistent models violating the energy budget balance of the atmosphere.
false
[ "non-consistent models", "non-equilibrium chemical abundances", "Chemical kinetics models", "non-equilibrium chemistry", "chemical equilibrium", "strong chemical disequilibrium", "Chemical", "fixed model input", "vertical mixing", "photochemistry", "consistent calculations", "hot Jupiters", "the non-consistent models", "P-T profiles", "the chemical kinetics calculation", "important consequences", "Jupiters", "nature", "Jupiter", "Simulated observations" ]
7.417861
14.641686
108
12500462
[ "Ciambur, Bogdan C." ]
2016PASA...33...62C
[ "Profiler - A Fast and Versatile New Program for Decomposing Galaxy Light Profiles" ]
43
[ "Centre for Astrophysics and Supercomputing, Swinburne University of Technology, Hawthorn, VIC 3122, Australia" ]
[ "2016ApJ...831..132G", "2017A&A...601A..20K", "2017A&A...602A.103W", "2017ApJ...840...68G", "2017ApJ...840...79S", "2017MNRAS.471.2321D", "2018ApJ...852..133S", "2018ApJ...869..113D", "2019A&A...632A.128B", "2019ApJ...873...85D", "2019ApJ...876..155S", "2019ApJ...887...10S", "2019MNRAS.487.1795S", "2019arXiv191002664R", "2019arXiv191007043B", "2020AN....341...10R", "2020ApJ...899...82B", "2020ApJ...900..178L", "2020ApJ...903...97S", "2020MNRAS.494.1751M", "2021ApJ...923..146G", "2021ApJ...923..246G", "2021IAUS..359...37D", "2021JApA...42...59S", "2021MNRAS.502.3085G", "2021MNRAS.503.2203C", "2021MNRAS.504.3831B", "2021MNRAS.508.1870S", "2021PASA...38...30D", "2021arXiv210208240P", "2022A&A...664A..92H", "2022ApJ...927...67S", "2022MNRAS.509.3938L", "2022MNRAS.514.3410H", "2022Natur.607..459B", "2023A&A...675A.105D", "2023ApJS..268...26O", "2023JPhCS2576a2013T", "2023MNRAS.518.4943D", "2023MNRAS.519.4651H", "2023MNRAS.520.1975G", "2024MNRAS.531.2223P", "2024arXiv240515427S" ]
[ "astronomy" ]
8
[ "galaxies: fundamental parameters", "galaxies: individual (NGC 2549", "NGC 3348", "Pox 52)", "galaxies: structure", "methods: data analysis", "Astrophysics - Instrumentation and Methods for Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1017/pasa.2016.60", "10.48550/arXiv.1607.08620" ]
1607
1607.08620_arXiv.txt
\label{sec:intro} Galaxies are complex structures assembled through a variety of physical processes which act at different stages of their life (e.g., gas accretion; star formation; disc formation, growth and buckling; bar formation and buckling etc.; as well as mergers and interactions with neighbouring galaxies). The result is a rich variety of galactic components in the observed galaxy population. Classifying galaxies based on these structures, in the optical and/or near-infrared bands has been and still is now common practice (e.g., \citealt{Jeans1919}, \citealt{Hubble1926}, \citealt{deVaucouleurs1959}, \citealt{Sandage1975}, \citealt{deVaucouleurs+1991}, \citealt{Abraham+2003}, \citealt{Buta+2015}, etc.). A quantitative structural classification requires a reliable method to separate out each structural component from the others that make up the galaxy. Moreover, individually analysing each constituent probes the specific physical or dynamical processes associated with it and thus provides insight into galaxy evolution. The common practice is to model one of the fundamental diagnostics of a galaxy's structure, namely its radial light (or surface brightness) profile (SBP), by decomposing it into a sum of analytical functions, with each function representing a single component (e.g., \citealt{Prieto+2001}, \citealt{Balcells+2003}, \citealt{Blanton+2003}, \citealt{Naab+2006}, \citealt{Graham&Worley2008}, etc.; the reader will find an insightful and comprehensive review of the long history of modelling galaxy light profiles in \citealt{Graham2013}). Galaxy SBPs are commonly extracted from images by fitting quasi-elliptical isophotes as a function of increasing distance (semi-major axis) from the photometric centre of the galaxy. In such schemes, the isophotes are free to change their axis ratio (ellipticity), position angle (PA) and shape (quantified through Fourier harmonics) with radius, which ensures that the models capture the galaxy light very well. A popular tool for this is the IRAF task {\sc Ellipse} (\citealt{Jedrzejewski1987}), which works well for galaxies whose isophotes display low-level deviations from pure ellipses (e.g., elliptical galaxies or disc galaxies viewed relatively face-on). For more complex isophotal structures however, (e.g., edge-on disc galaxies, X/peanut-shaped bulges, bars and barlenses) {\sc Ellipse} has been shown to fail and the newer IRAF task {\sc Isofit}\footnote{\url{https://github.com/BogdanCiambur/ISOFIT}} (\citealt{Ciambur2015}) is more appropriate. A somewhat different approach to performing galaxy decomposition is to directly fit the galaxy's (projected) light distribution in 2D (i.e., the galaxy image). Recent years have seen the advent and development of a number of programs dedicated to this purpose, notably {\sc GIM2D} (\citealt{Simard+2002}), {\sc Budda} (\citealt{deSouza+2004}), {\sc Galfit} (\citealt{Peng+2010}) and {\sc Imfit} (\citealt{Erwin2015}). In support for the 2D method, \cite{Erwin2015} has invoked several drawbacks of 1D profile modelling, namely that it is unclear which azimuthal direction to model (major axis, minor axis or other), that most of the data from the image is discarded, and that non-axisymmetric components (such as bars) can be misinterpreted as axisymmetric components, and their properties cannot be extracted from a 1D light profile. While these issues certainly apply when one extracts the SBP by taking a 1D cut from a galaxy image, all of these issues are resolved if the SBP is obtained from an isophotal analysis. In particular, fitting isophotes makes use of the entire image (so no data are discarded) and apart from the SBP itself, this process additionally provides information about the isophotes' ellipticities, position angles (PAs), and deviations from ellipticity (in the form of Fourier modes). All of this information is sufficient to completely reconstruct the galaxy image for even highly complex and non-axisymmetric isophote shapes (see \citealt{Ciambur2015} and Section \ref{sec:ex} of this paper). Having these extra isophote parameters allows one to obtain the SBP along any azimuthal direction, and identify and quantitatively study non-axisymmetric components such as bars or even peanut/X--shaped bulges (\citealt{Ciambur+2016}). It is therefore recommended to always use isophote tables rather than image cuts in 1D decompositions. Overall, both 1D and 2D decomposition techniques present benefits as well as disadvantages. The 2D image-modelling technique has the advantage that every pixel (except those deliberately masked out due to contaminating sources) in the image contributes directly to the fitting process, whereas in an isophotal (1D) analysis, pixels contribute in an azimuthal-average sense. Multicomponent systems with different photometric centres can also pose a problem for 1D SBPs, which assume a single centre for all components at $R=0$\footnote{This applies also to ring components, which have their brightest point at $R > 0$ along the 1D profile. This radial parameter represents the radius of the ring, while its centre is still assumed to be at $R = 0$.}, but can however be easily modelled in 2D. On the other hand, 2D codes suffer from the fact that each component has a single, fixed value for the ellipticity, PA, and Fourier moments (such as boxyness or discyness, and also higher orders), which can in some cases limit the method considerably. Triaxial ellipsoids viewed in projection can have radial gradients in their ellipticities and PAs (\citealt{Binney1978}, \citealt{Mihalas&Binney1981}), an effect captured in a 1D isophotal analysis (where both quantities can change with radius) but not in a 2D decomposition\footnote{Note, however, that the 2D code {\sc Imfit} can generate 2D images from line-of-sight integration of 3D luminosity density.}. There are notable examples in the literature where the 1D method has been preferred over the 2D technique. One such case is the decomposition of the {\sc Atlas$^{\rm 3D}$} (\citealt{Cappelari+2011}) sample{\footnote{\url{http://www-astro.physics.ox.ac.uk/atlas3d/}} of early-type galaxies, in \cite{Krajnovic+2013}. I point the reader to Sec. 2 and Appendix A of their paper, where they discuss both methods and test the performance of their preferred 1D method against a 2D analysis (with {\sc Galfit}). Another illuminating example is in \cite{Savorgnan&Graham2016}. They performed both 1D (with private code) and 2D (with {\sc Imfit}) decompositions of 72 galaxies, out of which 41 did not converge or did not give meaningful solutions in 2D, whereas only 9 could not be modelled in 1D. Sec. 4.1 in their paper also provides an insightful and practical comparison between 1D and 2D galaxy modelling techniques. The past few decades have seen a flurry of 2D image-fitting codes, whereas publicly available tools that focus on 1D decompositions are scarce. In this paper I present \prof, a freely-available code written in {\sc Python} and designed to provide a fast, flexible, user-friendly and accurate platform for performing structural decompositions of galaxy surface brightness profiles. The remainder of the paper is structured as follows. In Section \ref{sec:input} I describe the input data and information required by \prof~prior to the decomposition process. Section \ref{sec:model} is a concise review of typical galaxy components and the analytical functions employed to model them. Section \ref{sec:fitting} then details the fitting process, and Section \ref{sec:ex} provides three example applications, each illustrating different features of \prof: modelling a core-S\'ersic galaxy, the user-provided PSF vector feature and modelling a structurally complex edge-on galaxy with a spheroid, two nested bars and a truncated disc. Finally, I summarise and conclude with Section \ref{sec:cons}.
\label{sec:cons} I have introduced \prof, a flexible and user-friendly program coded in {\sc Python}, designed to model radial surface brightness profiles of galaxies. With an intuitive GUI, \prof~can model a wide range of galaxy components, such as spheroids (elliptical galaxies or the bulges of spiral or lenticular galaxies), face-on, inclined, edge-on and (anti)truncated discs, resolved or un-resolved nuclear point-sources, bars, rings, spiral arms, etc. with an arsenal of analytical functions routinely used in the field, such as the S\'ersic, core-S\'ersic, exponential, edge-on disc model, Gaussian, Moffat and Ferrers' functions. In addition, \prof~can employ a broken exponential model (relevant for disc truncations or antitruncations) and two 1D special cases of the 2D edge-on disc model, namely along the major axis (in the disc plane) and along the minor axis (perpendicular to the disc plane). \prof~is optimised to analyse isophote tables generated by the IRAF tasks {\sc Ellipse} and {\sc Isofit} but can also analyse two-column tables of radius and surface brightness. After reaching the best-fitting solution, the corresponding model parameters are returned. The major and equivalent axis profiles can both be analysed, and for the latter profile, each component's total magnitude is additionally returned. The model convolution with the PSF is performed in 2D, with an FFT-based scheme. This allows for elliptical models, and additionally ensures that the convolution conserves the model's total flux (as a 1D convolution of the model profile with the PSF profile does not). Further, \prof~allows for a choice between Gaussian, Moffat or a user-provided data vector for the PSF (a table of $R$ and $I(R)$ values). All of the possible PSF choices can also be used as point-source components in the model. \prof~is freely available from the following URL: \url{https://github.com/BogdanCiambur/PROFILER}.
16
7
1607.08620
I introduce Profiler, a user-friendly program designed to analyse the radial surface brightness profiles of galaxies. With an intuitive graphical user interface, Profiler can accurately model galaxies of a broad range of morphological types, with various parametric functions routinely employed in the field (Sérsic, core-Sérsic, exponential, Gaussian, Moffat, and Ferrers). In addition to these, Profiler can employ the broken exponential model for disc truncations or anti-truncations, and two special cases of the edge-on disc model: along the disc's major or minor axis. The convolution of (circular or elliptical) models with the point spread function is performed in 2D, and offers a choice between Gaussian, Moffat or a user-provided profile for the point spread function. Profiler is optimised to work with galaxy light profiles obtained from isophotal measurements, which allow for radial gradients in the geometric parameters of the isophotes, and are thus often better at capturing the total light than 2D image-fitting programs. Additionally, the 1D approach is generally less computationally expensive and more stable. I demonstrate Profiler's features by decomposing three case-study galaxies: the cored elliptical galaxy NGC 3348, the nucleated dwarf Seyfert I galaxy Pox 52, and NGC 2549, a double-barred galaxy with an edge-on, truncated disc.
false
[ "galaxy light profiles", "disc truncations", "galaxies", "spread function", "various parametric functions", "radial gradients", "2D", "Moffat", "the radial surface brightness profiles", "Gaussian", "Sérsic", "2D image-fitting programs", "isophotal measurements", "truncations", "NGC", "morphological types", "the cored elliptical galaxy", "Ferrers", "exponential", "-" ]
11.364939
5.003278
-1
1936667
[ "Díaz, Mario C.", "Beroiz, Martín", "Peñuela, Tania", "Macri, Lucas M.", "Oelkers, Ryan J.", "Yuan, Wenlong", "García Lambas, Diego", "Cabral, Juan", "Colazo, Carlos", "Domínguez, Mariano", "Sánchez, Bruno", "Gurovich, Sebastián", "Lares, Marcelo", "Schneiter, Matías", "Graña, Darío", "Renzi, Víctor", "Rodriguez, Horacio", "Starck, Manuel", "Vrech, Rubén", "Artola, Rodolfo", "Chiavassa Ferreyra, Antonio", "Girardini, Carla", "Quiñones, Cecilia", "Tapia, Luis", "Tornatore, Marina", "Marshall, Jennifer L.", "DePoy, Darren L.", "Branchesi, Marica", "Brocato, Enzo", "Padilla, Nelson", "Pereyra, Nicolas A.", "Mukherjee, Soma", "Benacquista, Matthew", "Key, Joey" ]
2016ApJ...828L..16D
[ "GW150914: First Search for the Electromagnetic Counterpart of a Gravitational-wave Event by the TOROS Collaboration" ]
15
[ "Center for Gravitational Wave Astronomy, University of Texas Rio Grande Valley, Brownsville, TX, USA", "Center for Gravitational Wave Astronomy, University of Texas Rio Grande Valley, Brownsville, TX, USA; University of Texas at San Antonio, San Antonio, TX, USA", "Center for Gravitational Wave Astronomy, University of Texas Rio Grande Valley, Brownsville, TX, USA; Ludwig Maximilian Universität Munich, Faculty of Physics, Munich, Germany", "Mitchell Institute for Fundamental Physics &amp; Astronomy, Department of Physics &amp; Astronomy, Texas A&amp;M University, College Station, TX, USA", "Mitchell Institute for Fundamental Physics &amp; Astronomy, Department of Physics &amp; Astronomy, Texas A&amp;M University, College Station, TX, USA;", "Mitchell Institute for Fundamental Physics &amp; Astronomy, Department of Physics &amp; Astronomy, Texas A&amp;M University, College Station, TX, USA", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina; Ministerio de Educación de la Provincia de Córdoba, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Universidad Nacional de Córdoba, IATE, Córdoba, Argentina", "Mitchell Institute for Fundamental Physics &amp; Astronomy, Department of Physics &amp; Astronomy, Texas A&amp;M University, College Station, TX, USA", "Mitchell Institute for Fundamental Physics &amp; Astronomy, Department of Physics &amp; Astronomy, Texas A&amp;M University, College Station, TX, USA", "Università degli studi di Urbino, Urbino, Italy", "INAF—Osservatorio Astronomico di Roma, Monte Porzio Catone, Italy", "Instituto de Astrofísica, Pontificia Universidad Católica de Chile, Santiago, Chile", "Center for Gravitational Wave Astronomy, University of Texas Rio Grande Valley, Brownsville, TX, USA", "Center for Gravitational Wave Astronomy, University of Texas Rio Grande Valley, Brownsville, TX, USA", "Center for Gravitational Wave Astronomy, University of Texas Rio Grande Valley, Brownsville, TX, USA", "Center for Gravitational Wave Astronomy, University of Texas Rio Grande Valley, Brownsville, TX, USA" ]
[ "2016arXiv160909517A", "2017ApJ...846...62S", "2017ApJ...848L..29D", "2017ApJ...850...21A", "2017GReGr..49...83M", "2017MNRAS.466L..78V", "2018AJ....155...39O", "2018AJ....156..132O", "2018ApJ...857...81G", "2018LRR....21....3A", "2019A&C....2800284S", "2019ApJ...886...73N", "2019JPhCS1390a2088A", "2019MNRAS.484.3219M", "2022AJ....163...95S" ]
[ "astronomy" ]
2
[ "gravitational waves", "techniques: image processing", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "1987PASP...99..191S", "1991ApJ...380L..17P", "1996A&AS..117..393B", "1996ApJ...473..356P", "1998ApJ...503..325A", "1998ApJ...507L..59L", "2000A&AS..144..363A", "2002ApJ...572..407B", "2003AAS...203.9601M", "2003AJ....125..984M", "2004AAS...20510818G", "2006AJ....131.1163S", "2008MNRAS.386L..77B", "2008PASP..120..449M", "2009BAAA...52..285R", "2009PhDT.......235L", "2010PhRvD..82j2002N", "2011CQGra..28h5016W", "2011MNRAS.413.2004C", "2011Natur.478...82N", "2012A&A...541A.155A", "2012A&A...548A..65T", "2012ApJ...746...48M", "2013ApJ...775...18B", "2013ApJ...776...18F", "2014ARA&A..52...43B", "2014ApJ...784....8H", "2015AJ....149...50O", "2015ApJ...800...81E", "2015ApJ...814...25C", "2015CQGra..32b4001A", "2015CQGra..32g4001L", "2015GCN.18330....1S", "2016A&A...594A..13P", "2016AJ....151..166O", "2016ApJ...823L..33S", "2016ApJ...823L..34A", "2016ApJ...824L..24K", "2016ApJS..225....8A", "2016LRR....19....1A", "2016MNRAS.462.4094S", "2016PhRvD..93d2004K", "2016PhRvL.116f1102A", "2017ApJ...835...64G", "2017MNRAS.465.3656L", "2017PhRvD..95j4046L" ]
[ "10.3847/2041-8205/828/2/L16", "10.48550/arXiv.1607.07850" ]
1607
1607.07850_arXiv.txt
\setcounter{footnote}{10} The network of advanced ground-based gravitational wave (GW) interferometers constituted by the LIGO observatories \citep{LSC2015}, which started operations in September 2015 and by the VIRGO observatory \citep{Acernese2015}, which will join before the end of 2016, were designed to be capable of detecting GWs emitted by the mergers of neutron stars and/or black holes in binary systems out to distances of hundreds of Mpc \citep[see][and references therein]{Abbott2016b}. In anticipation of the operation of this network, on 2013 June 6 the LIGO-VIRGO collaboration (LVC) issued a worldwide call\footnote{http://www.ligo.org/scientists/GWEMalerts.php} to astronomers to participate in multi-messenger observations of astrophysical events recorded by the GW detectors, using a wide range of telescopes and instruments of ``mainstream'' astronomy. \ \par Initially, triggers will be shared promptly only with astronomy partners who have signed a Memorandum of Understanding (MoU) with LVC involving an agreement on deliverables, publication policies, confidentiality, and reporting. It is expected that if the mergers of compact objects contain at least one neutron star, electromagnetic (EM) radiation will be emitted during the event. This EM counterpart, originating in the ejecta and its interaction with the surrounding environment could range from very short duration gamma-ray bursts to longer-duration emission at optical, near infrared (kilonova and short GRB afterglows) and radio wavelengths \citep[e.g.,][]{Li1998,Nakar2011,Metzger2012,Barnes2013,Berger2014,Cowperthwaite2015}. Simultaneous detection of the event by GW and EM observatories could provide a more integrated astrophysical interpretation of the event and would be instrumental in producing better estimates for the distance and energy scales of the event. Motivated to participate in these observations, we formed a collaboration named ``Transient Optical Robotic Observatory of the South'' \citep[TOROS;][]{Benacquista2014} which seeks to deploy a wide-field optical telescope on Cord\'on Mac\'on in the Atacama Plateau of northwestern Argentina \citep{Renzi2009,Tremblin2012}. The collaboration planned to utilize other resources independently of the construction of this facility. On 2014 April 5, the TOROS collaboration signed a Memorandum of Understanding with LVC and participated during the first scientific run of the GW interferometers from September 2015 through January 2016. Two facilities were available to TOROS during this campaign: a Schmidt-Cassegrain 0.4-m telescope at Cord\'on Mac\'on and a 1.5-m telescope at Estaci\'on Astrof\'{\i}sica Bosque Alegre (EABA) in C\'ordoba, Argentina. On 2015 September 14 at 09:50:45 UT, the two USA-based detectors of the Advanced LIGO interferometer network detected a high-significance candidate GW event designated GW150914 \citep{Abbott2016a}. This unexpected detection --- observed four days before the first scientific run of the detectors was scheduled to start --- constituted the first detection of the merger of a binary black hole (BBH) system and the first direct detection of gravitational waves. Due to the unexpected timing of the event, LVC provided spatial location information two days later, in the form of probability sky maps via a private GCN circular \citep[GCN\#18330]{Singer2015}. TOROS was one of 25 teams that participated in the search for an electromagnetic counterpart in the southern hemisphere. We report here on the optical follow up of this event by the TOROS collaboration during 2015 September 16-17 using the 1.5-m EABA telescope (the smaller telescope at Mac\'on was not operational at the time). This paper is organized as follows: \S\ref{observations} discusses target selection and observations; \S\ref{data-analysis} describes the data reduction, image differencing algorithms and the bogus/real classification; \S\ref{results} presents our results and \S\ref{conclusions} summarizes our findings. Throughout this paper, we express magnitudes in the AB system and adopt a $\Lambda$CDM cosmology based on results from the Planck mission \citep{Planck2015}.
\label{conclusions} The TOROS collaboration conducted a prompt search for the electromagnetic counterpart of the first gravitational-wave event reported by LIGO using the 1.5-m telescope of Estaci\'on Astrof\'{\i}sica Bosque Alegre (EABA) in C\'ordoba, Argentina. Our search spanned two nights, during which we targeted 21 fields containing 14 nearby ($D<60$~Mpc) galaxies with high probabilities of hosting the event. We covered $0.62\sq^{\circ}$ and reached a $5\sigma$ limiting AB magnitude of $r=21.7$. We used a combination of difference-imaging techniques and machine-learning procedures to detect and classify potential transients. No {\it bona fide} events were found, a result that is consistent with the low probability of detecting stellar or extragalactic variability given our temporal and areal coverage, and with the later classification of the GW event as a merger of two stellar-mass black holes. Our host-galaxy ranking approach serves as a complementary strategy to the wide-field surveys for these transients, such as those conducted by the Dark Energy Survey \citep{Annis2016,Soares2016}, the intermediate Palomar Transient Factory \citep{Kasliwal2016}, MASTER \citep{2016arXiv160501607L}, Pan-STARSS \citep{2016arXiv160204156S}, and the VLT Survey Telescope (E. Brocato et al, in preparation). Given the incompleteness of local galaxy catalogs, the rapid dissemination of possible counterpart candidates by the wide-field surveys would enable detailed photometric coverage to be contributed by many modest-aperture, narrow-field telescopes throughout the world. Additionally, unfiltered CCD observations may be desirable at this stage given the large uncertainties in the possible colors of these counterparts.
16
7
1607.07850
We present the results of the optical follow-up conducted by the TOROS collaboration of the first gravitational-wave event GW150914. We conducted unfiltered CCD observations (0.35-1 μm) with the 1.5 m telescope at Bosque Alegre starting ∼2.5 days after the alarm. Given our limited field of view (∼100 arcmin<SUP>2</SUP>), we targeted 14 nearby galaxies that were observable from the site and were located within the area of higher localization probability. We analyzed the observations using two independent implementations of difference-imaging algorithms, followed by a Random-Forest-based algorithm to discriminate between real and bogus transients. We did not find any bona fide transient event in the surveyed area down to a 5σ limiting magnitude of r = 21.7 mag (AB). Our result is consistent with the LIGO detection of a binary black hole merger, for which no electromagnetic counterparts are expected, and with the expected rates of other astrophysical transients.
false
[ "other astrophysical transients", "higher localization probability", "real and bogus transients", "Bosque Alegre", "r", "days", "magnitude", "unfiltered CCD observations", "5σ", "the first gravitational-wave event GW150914", "a binary black hole merger", "first", "any bona fide transient event", "the expected rates", "the surveyed area", "∼100", "difference-imaging algorithms", "view", "Random-Forest", "(AB" ]
7.196929
2.079722
62
3182220
[ "Bento, J.", "Schmidt, B.", "Hartman, J. D.", "Bakos, G. Á.", "Ciceri, S.", "Brahm, R.", "Bayliss, D.", "Espinoza, N.", "Zhou, G.", "Rabus, M.", "Bhatti, W.", "Penev, K.", "Csubry, Z.", "Jordán, A.", "Mancini, L.", "Henning, T.", "de Val-Borro, M.", "Tinney, C. G.", "Wright, D. J.", "Durkan, S.", "Suc, V.", "Noyes, R.", "Lázár, J.", "Papp, I.", "Sári, P." ]
2017MNRAS.468..835B
[ "HATS-22b, HATS-23b and HATS-24b: three new transiting super-Jupiters from the HATSouth project" ]
15
[ "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, Australian National University, Cotter Road, Weston, ACT 2611, Australia", "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, Australian National University, Cotter Road, Weston, ACT 2611, Australia", "Department of Astrophysical Sciences, Princeton University, 4 Ivy Ln., Princeton, NJ 08544, USA", "Department of Astrophysical Sciences, Princeton University, 4 Ivy Ln., Princeton, NJ 08544, USA", "Max Plank Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany; Department of Astronomy, Stockholm University, SE-106 91 Stockholm, Sweden", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, Av. Vicuña Mackenna 4860, 7820436 Macul, Santiago, Chile", "Observatoire Astronomique de l'Université de Genéve, 51 ch. des Maillettes, CH-1290 Versoix, Switzerland", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, Av. Vicuña Mackenna 4860, 7820436 Macul, Santiago, Chile", "Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA", "Max Plank Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany; Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, Av. Vicuña Mackenna 4860, 7820436 Macul, Santiago, Chile", "Department of Astrophysical Sciences, Princeton University, 4 Ivy Ln., Princeton, NJ 08544, USA", "Department of Astrophysical Sciences, Princeton University, 4 Ivy Ln., Princeton, NJ 08544, USA", "Department of Astrophysical Sciences, Princeton University, 4 Ivy Ln., Princeton, NJ 08544, USA", "Max Plank Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany; Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, Av. Vicuña Mackenna 4860, 7820436 Macul, Santiago, Chile", "Max Plank Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany", "Max Plank Institute for Astronomy, Königstuhl 17, D-69117 Heidelberg, Germany", "Department of Astrophysical Sciences, Princeton University, 4 Ivy Ln., Princeton, NJ 08544, USA", "Australian Centre for Astrobiology, School of Physics, University of New South Wales, NSW 2052, Australia; Exoplanetary Science at UNSW, School of Physics, University of New South Wales, NSW 2052, Australia", "Australian Centre for Astrobiology, School of Physics, University of New South Wales, NSW 2052, Australia; Exoplanetary Science at UNSW, School of Physics, University of New South Wales, NSW 2052, Australia", "Astrophysics Research Centre, School of Mathematics &amp; Physics, Queen's University, Belfast BT7 1NN, UK", "Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica de Chile, Av. Vicuña Mackenna 4860, 7820436 Macul, Santiago, Chile", "Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA", "Hungarian Astronomical Association, Budapest, 1461 Hungary", "Hungarian Astronomical Association, Budapest, 1461 Hungary", "Hungarian Astronomical Association, Budapest, 1461 Hungary" ]
[ "2016AJ....152..182H", "2017MNRAS.467..971B", "2018AJ....155..119B", "2018exha.book.....P", "2019AJ....158...59S", "2019RNAAS...3...35O", "2020AJ....160..155W", "2021AJ....161..250W", "2021CoSka..51...68G", "2022ApJS..258...40K", "2022ApJS..259...62I", "2022ExA....53..547K", "2022MNRAS.513..102C", "2023ApJS..265....4K", "2023MNRAS.525..876M" ]
[ "astronomy" ]
31
[ "techniques: photometric", "techniques: spectroscopic", "stars: individual: HATS-22", "stars: individual: GSC 6664-00373", "stars: individual: HATS-23", "stars: individual: GSC 8382-01464", "stars: individual: HATS-24", "stars: individual: GSC 9054-00129", "planetary systems", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1952Obs....72..199S", "1962AJ.....67..591K", "1989ApJ...345..245C", "1995Icar..114..217L", "1995Sci...267..360B", "1996Sci...274..954R", "1997ApJ...478..367M", "1998SPIE.3355..844K", "2000A&AS..141..371G", "2000A&AS..143...23O", "2000ApJ...537..916S", "2000Icar..143....2B", "2001ApJS..136..417Y", "2001Msngr.105....1Q", "2002A&A...391..369K", "2002ApJ...580L.171M", "2002PASP..114..144F", "2003ApJ...589..605W", "2003Msngr.114...20M", "2004A&A...428.1001C", "2005A&A...434..343A", "2005ApJ...619..570M", "2005MNRAS.356..557K", "2007Ap&SS.310..255D", "2007ApJ...655..541M", "2007ApJ...659.1661F", "2007ApJ...671..861H", "2008A&A...486..951G", "2008ApJ...673..526K", "2008ApJ...680.1450P", "2008ApJ...686..621F", "2008MNRAS.384.1242R", "2008MNRAS.385..109P", "2009A&A...496..577Z", "2009ApJ...705.1206C", "2009ApJ...707..167S", "2009Msngr.137...14H", "2010ApJ...710.1724B", "2010MNRAS.403..151C", "2011AJ....141...30C", "2011ApJ...735..109W", "2011ApJ...739...31L", "2011ApJS..197...11D", "2011MNRAS.414.1278P", "2011MNRAS.416.1443S", "2012A&A...543L...5G", "2012AJ....144..139H", "2013A&A...552A.120S", "2013A&A...556A.150S", "2013AJ....145....5P", "2013AJ....146..113B", "2013ApJ...774..118Z", "2013ApJ...774L...9A", "2013ApJ...778..184J", "2013PASP..125..154B", "2013PASP..125.1031B", "2013arXiv1301.3156W", "2013cctp.book..367M", "2014A&A...562A.140P", "2014AJ....148...29J", "2014ARA&A..52..171O", "2014MNRAS.437.1511B", "2014MNRAS.445.2746Z", "2015AJ....150...49B", "2015AJ....150..197H", "2015ApJ...805...75P", "2015ApJ...812...94D", "2015ApJ...814L..24L", "2016AJ....151...89B", "2016AJ....152..161D", "2016AJ....152..174A", "2016ApJ...825...62S", "2016ApJ...827....8N", "2016ApJ...829..114B", "2016MNRAS.460.3376Z", "2016Natur.534..662D", "2017MNRAS.467..971B" ]
[ "10.1093/mnras/stx500", "10.48550/arXiv.1607.00688" ]
1607
1607.00688_arXiv.txt
\label{sec:introduction} Transiting planets are the key towards understanding the structure and composition of planetary systems. The breadth of system parameters that can be determined from the discovery data sets and follow-up studies surpasses any other detection method, the most important being the mass and radius, yielding an estimate of the bulk density. Moreover, these planets are amenable to transmission studies during transit \citep[e.g.][]{bento:2014,jordan:2013,marley:2013,seager:2000, pont:2008,sing:2011}, a direct measurement of the planet's day-side emission as an estimate of the surface temperature during secondary eclipse \citep{zhou:2013,zhou:2014:sec,knutson:2008,desert:2011,croll:2011}, and other properties \citep{louden:2015,zhou:2016,collier:2010, hartman:2015, gandolfi:2012}. In particular, the hundreds of \hjs (broadly Jupiter mass planets orbiting close to their host stars on less that $\sim 10$ day periods) found to date have challenged planetary formation theories and structure models. Despite an early suggestion of the possibility that such planets may exist by \citet{struve:1952}, explaining their existence is not trivial as they are not generally expected to form in-situ \citep{boss:1995, lissauer:1995,bodenheimer:2000}, with the migration potentially taking place in the very early early stages of formation \citep{donati:2016}. Recent work suggests that there is a potential mechanism that can form such planets in-situ \citep{batygin:2016}, but the general consensus is that these planets are formed at large separations and migrate inwards to their current positions, and several possible mechanisms have been suggested for this process. A disk migration scenario has been proposed in which the orbiting planet exchanges angular momentum with the protoplanetary disk and loses orbital momentum, thereby starting out at large separations and making its way in to close to the host star \citep[e.g.][and references therein]{alibert:2005,chambers:2009,rice:2008}. Alternatively, interactions with other bodies in the system can cause scattering/ejection events and the planet in question is forced into an eccentric orbit that brings it closer to the host star \citep[e.g.][]{rasio:1996,ford:2008}. Tidal interactions are then thought to circularise the orbit resulting in close-in planets. Other processes suggested include Kozai migration, first proposed by \cite{wu:2003}, which states that a highly inclined stellar companion can induce Kozai oscillations \citep{kozai:1962} in the planet and excite it to progressively higher eccentricity. We note, however, that \citet{ngo:2016} suggest that only a small fraction ($<$20\%) of \hj host stars have stellar companions capable of inducing such oscillations. Very recent works by \cite{petrovich:2015} and \cite{wu:2011} suggest that {\em secular migrations} may occur due to interactions between two-or-more well-spaced, eccentric planets, which can cause one of them to become very eccentric on long timescales, leading to both enhanced eccentricity and tidal dissipation over larger timescales \citep{lithwick:2011}. However, recent results show that an understanding of planet formation and migration has not been achieved yet. \citet{antonini:2016} suggest that perhaps \hjs with outer companions are unlikely to have migrated through high-eccentricity processes due to the instability of their orbits, while \citet{schlaufman:2016} find that warm Jupiters are no more likely to have wide orbit planetary companions than those in longer orbits, which is at odds with an eccentric migration scenario. The possibility that there is a mass dependence in the question of planet migration and eccentricity is supported by evidence that higher mass planets tend to show higher eccentricity than those less massive than 2$\mjup$ \citep{mazeh:1997,marcy:2005,southworth:2009}. Moreover, it seems that planets at higher orbital separation/period also have a higher tendency to show non-zero eccentricities \citep{pont:2011} versus close-in planets. This raises questions such as: are high-mass planets more susceptible to retain large eccentricities on longer timescales? And, if so, is this an indication that planet-planet scattering, predicted to generate high eccentric orbits, is likely to be the main migration mechanism for planetary systems? Is the structure and evolution of high and low-mass planets fundamentally different? Are \hj structures fundamentally affected by extreme cases of inwards migration and current irradiation levels? The answer lies on a larger sample and better understanding of the composition of these planets. In this paper we report the discovery of three new transiting super Jupiters with masses higher than 1.4\mjup\, from the HATSouth survey: \hatcurb{22}, \hatcurb{23} and \hatcurb{24}. These planets add to the list of of well-characterised massive \hjs which collectively pose a challenge to models of planetary formation and migration. In Section \ref{sec:obs} we describe the photometric and spectroscopic observations undertaken for all three targets in the pursuit of determining their planetary nature. Section \ref{sec:analysis} contains a description of the global data analysis and presents the modelled stellar and planetary parameters. We also describe the methods employed to reject false positive scenarios. Our findings are finally discussed in Section \ref{sec:discussion}.
\label{sec:discussion} \begin{figure} { \centering \includegraphics[width={\linewidth}]{massradius.eps} } \caption{ Mass-radius relation for \hjs, defined as those planets with masses higher than $0.5 \mjup$ and periods shorter than 10 days. We show theoretical models for planet structures from \cite{fortney:2007} for each of the three planets announced in this paper for both no core (solid lines) and $100 \mearth$ core (dashed lines) scenarios. The new HATSouth planets are indicated. 4.5 Gy 0.1 AU models are shown for comparison with \hatcurb{22} (yellow lines). The 4.5Gy 0.02 AU models (red lines) are used as equivalent examples for \hatcurb{23} and the models for \hatcurb{24} (1 Gy 0.02 AU) is shown in green lines. } \label{fig:massradius} \end{figure} \begin{figure} { \centering \includegraphics[width={\linewidth}]{arsmratio.eps} } \caption{ Planet to star mass ratio as a function of the ratio between the orbital separation and the stellar radius for known \hjs. We have labelled our new discoveries in the plot, which are marked in filled red symbols. The dot sizes correspond to the planets' measured eccentricity ranging from zero (smallest size) up to a value of 0.562. } \label{fig:arsmratio} \end{figure} In this work we report the discovery of three new moderately high-mass \hjs by the HATSouth survey: \hatcurb{22}, \hatcurb{23} and \hatcurb{24}. These planets add to the growing numbers of known \hjs and provide vital additional insights into the formation and distribution of short-period massive planets. In Figures \ref{fig:massradius} and \ref{fig:arsmratio} we show these discoveries in the context of known planets with masses higher than $0.5 \mjup$ and less than 10 day orbital periods \footnote{Previously known planets taken from the NASA Exoplanet Archive at http://exoplanetarchive.ipac.caltech.edu/ on 02/02/2017}. \subsection{Mass-radius relation} \label{sec:massradius} In Figure \ref{fig:massradius} we plot a mass-radius relation for known \hjs, highlighting our three new discovered planets. Additionally, this plot contains a selection of predicted mass-radius relations from \cite{fortney:2007} that are closest to the conditions of each of these planets. For \hatcurb{23} and \hatcurb{24} (red and green curves respectively) we have selected those models that most closely matched the determined age and orbital separation of the planets (4.5 and 1 Gy respectively at 0.02 au separation) and with orbital periods of less than 10 days. These models assume a solar analogue host star and solar luminosity. For \hatcurb{22}, however, due to the lower luminosity of the host star, we have selected the curve that most closely matches the expected equilibrium temperature $T_{eq}$ for this planet, a 4.5 Gy model at 0.1 au, which is equivalent to a $T_{eq}$ of approximately 875K, thereby matching the model to the planet's conditions. In all cases, we plot the two extreme limits of planetary core mass: no core (solid line) and a $100 \mearth$ core model (dashed line). Despite the fact that none of the target's radii fall within their equivalent predicted mass-radius range from the model curves, \hatcurb{22} is consistent with a high-mass ($100 \mearth$) core at a 95\% confidence level. This places this planet among the twenty highest bulk density \hjs known to date, defined as those with masses higher than $0.5 \mjup$ and less than 10 day orbital periods. While we are not able to confidently make any conclusions regarding the structure of grazing transiting planet \hatcurb{23}, we note that even assuming the lower limit on the radius of this planet at $1.31 \mjup$ it is still more likely that this is an example of a low core mass \hj. Perhaps the most interesting case in this context is that of \hatcurb{24}. The radius of this planet is estimated more than $3 \sigma$ above the model line for a pure helium-hydrogen planet at approximately this age and orbital separation. The short period and young age of the host star leads to a high equilibrium temperature ($T_{eq} = 2067 \pm 39 K$), which puts it in a similar regime to other inflated \hjs (e.g. WASP-71b \citep{smith:2013} and HATS-35b \citep{valborro:2016}). As noted by \cite{marley:2007} and \cite{fortney:2007}, the physical properties of giant planets at young ages are uncertain, and the model shown does not include considerations on the formation mechanism. The somewhat inflated radius of this planet may be due to a combination of mechanisms discussed in Section \ref{sec:introduction}. Additionally, the apparent higher radius may be due to an ongoing evaporation of the upper layers of the planet's atmosphere as current models are unable to explain the radius of \hatcurb{24}. This factor, coupled with the fact that the host star is moderately bright ($V = 12.830 \pm 0.010$) and the large transit depth signal, makes this a good target for further wavelength dependent transmission studies with large-class telescopes. Assuming a hydrogen dominated atmosphere, the scale height $H$ for this planet is estimated at $314.8 \pm 32.8$ km, a value comparable to other inflated hot-Jupiters. Therefore, a transmission spectroscopic signal equivalent to 5 scale heights would be detectable at a level of $\sim$250ppm, well within the capability of current and future facilities. \subsection{Eccentricity and tidal circularisation} A comparative study between the three announced planets and previously known \hjs is useful as the orbital eccentricity is an important parameter thought to be highly related to planetary migration both in disk interaction and planet scattering scenarios. As discussed in Section \ref{sec:introduction}, eccentric orbits are found preferentially for higher mass planets \citep[e.g.][]{southworth:2009} and a higher the orbital separation seems to correspond with non-zero eccentricity \citep{pont:2011}. This supports the fact that a high mass planet at a large separation should require more time for tidal circularisation. Our analysis (discussed in Section \ref{sec:globmod}) finds \hatcurb{22} to be the only planet of the three with a likely non-zero eccentricity of $e = \hatcurRVecceneccen{22}$. The distinction between \hatcurb{22} and the two other planets discovered in this work is made more evident in Figure \ref{fig:arsmratio}, in which we have plotted the planet to stellar mass ratio as a function of the ratio between the orbital separation and the stellar radius ($a/R_*$). These parameters are the dominant factors in determining the tidal circularisation timescale of planetary orbits \citep{duffell:2015,ogilvie:2014} in which the mass ratio is proportional to the circularisation timescale. This is due to the fact that more massive planets are more likely to carve a large gap in the protoplanetary disk, leading to less efficient circularisation due to disk interaction. Thus, these planets should remain in moderate non-zero eccentric orbits for a longer time. In Figure \ref{fig:arsmratio} the dot sizes correspond to the eccentricity values, ranging from zero to 0.562, where the concentration of low or zero eccentricity planets, can be found for low values of both plotted parameters. \hatcurb{22} can be seen occupying a region of parameter space clearly distinguished from that of \hatcurb{23} and \hatcurb{24}, suggesting that the eccentric orbit of \hatcurb{22} may indeed be a result of insufficient time for full tidal circularisation. However, several examples can still be found in this figure which are inconsistent with this picture. The case of CoRoT-27b \citep{parviainen:2014} is of a non-eccentric orbit planet that can be seen as the point with a mass ratio just under 0.01. Despite the moderate age of this system ($4.21 \pm 2.72$ Gyr), a planet with such a high mass ratio is still predicted to be found in an eccetric orbit. On the other hand, several cases can also be found with low values of both parameters plotted in Figure \ref{fig:arsmratio} that still presently show detectable orbital eccentricity, suggesting that there are other mechanisms, besides tidal circularisation, determining the eccentricity distribution of exoplanets. \ifthenelse{\boolean{emulateapj}}{ \begin{deluxetable*}{lccc} }{ \begin{deluxetable}{lccc} } \tabletypesize{\scriptsize} \tablecaption{Orbital and planetary parameters for \hatcurb{22}, \hatcurb{23} and \hatcurb{24}\label{tab:planetparam}} \tablehead{ \multicolumn{1}{c}{} & \multicolumn{1}{c}{\bf HATS-22b} & \multicolumn{1}{c}{\bf HATS-23b} & \multicolumn{1}{c}{\bf HATS-24b} \\ \multicolumn{1}{c}{~~~~~~~~~~~~~~~Parameter~~~~~~~~~~~~~~~} & \multicolumn{1}{c}{Value} & \multicolumn{1}{c}{Value} & \multicolumn{1}{c}{Value} } \startdata \noalign{\vskip -3pt} \sidehead{\Lc{} parameters} ~~~$P$ (days) \dotfill & $\hatcurLCPeccen{22}$ & $\hatcurLCP{23}$ & $\hatcurLCP{24}$ \\ ~~~$T_c$ (${\rm BJD}$) \tablenotemark{a} \dotfill & $\hatcurLCTeccen{22}$ & $\hatcurLCT{23}$ & $\hatcurLCT{24}$ \\ ~~~$T_{14}$ (days) \tablenotemark{a} \dotfill & $\hatcurLCdureccen{22}$ & $\hatcurLCdur{23}$ & $\hatcurLCdur{24}$ \\ ~~~$T_{12} = T_{34}$ (days) \tablenotemark{a} \dotfill & $\hatcurLCingdureccen{22}$ & $\hatcurLCingdur{23}$ & $\hatcurLCingdur{24}$ \\ ~~~$\arstar$ \dotfill & $\hatcurPPareccen{22}$ & $\hatcurPPar{23}$ & $\hatcurPPar{24}$ \\ ~~~$\zrstar$ \tablenotemark{b} \dotfill & $\hatcurLCzetaeccen{22}$\phn & $\hatcurLCzeta{23}$\phn & $\hatcurLCzeta{24}$\phn \\ ~~~$\rpl/\rstar$ \dotfill & $\hatcurLCrprstareccen{22}$ & $\hatcurLCrprstar{23}$ & $\hatcurLCrprstar{24}$ \\ ~~~$b^2$ \dotfill & $\hatcurLCbsqeccen{22}$ & $\hatcurLCbsq{23}$ & $\hatcurLCbsq{24}$ \\ ~~~$b^2$ lower limit \tablenotemark{c} \dotfill & $\cdots$ & $\hatcurLCbsqsiglowerlim{23}$ & $\cdots$ \\ ~~~$b \equiv a \cos i/\rstar$ \dotfill & $\hatcurLCimpeccen{22}$ & $\hatcurLCimp{23}$ & $\hatcurLCimp{24}$ \\ ~~~$b$ lower limit \tablenotemark{c} \dotfill & $\cdots$ & $\hatcurLCimpsiglowerlim{23}$ & $\cdots$ \\ ~~~$i$ (deg) \dotfill & $\hatcurPPieccen{22}$\phn & $\hatcurPPi{23}$\phn & $\hatcurPPi{24}$\phn \\ ~~~$i$ upper limit (deg) \tablenotemark{c} \dotfill & $\cdots$\phn & $\hatcurPPitwosigupperlim{23}$\phn & $\cdots$\phn \\ \sidehead{HATSouth blend factors \tablenotemark{d}} ~~~Blend factor \dotfill & $\hatcurLCiblend{22}$ & $\hatcurLCiblend{23}$ & $\hatcurLCiblend{24}$ \\ \sidehead{Limb-darkening coefficients \tablenotemark{e}} ~~~$c_1,r$ \dotfill & $\hatcurLBireccen{22}$ & $\hatcurLBir{23}$ & $\hatcurLBir{24}$ \\ ~~~$c_2,r$ \dotfill & $\hatcurLBiireccen{22}$ & $\hatcurLBiir{23}$ & $\hatcurLBiir{24}$ \\ ~~~$c_1,i$ \dotfill & $\hatcurLBiieccen{22}$ & $\hatcurLBii{23}$ & $\hatcurLBii{24}$ \\ ~~~$c_2,i$ \dotfill & $\hatcurLBiiieccen{22}$ & $\hatcurLBiii{23}$ & $\hatcurLBiii{24}$ \\ ~~~$c_1,z$ \dotfill & $\cdots$ & $\hatcurLBiz{23}$ & $\cdots$ \\ ~~~$c_2,z$ \dotfill & $\cdots$ & $\hatcurLBiiz{23}$ & $\cdots$ \\ \sidehead{radial velocity parameters} ~~~$K$ (\ms) \dotfill & $\hatcurRVKeccen{22}$\phn\phn & $\hatcurRVK{23}$\phn\phn & $\hatcurRVK{24}$\phn\phn \\ ~~~$e$ \tablenotemark{f} \dotfill & $\hatcurRVecceneccen{22}$ & $\hatcurRVeccentwosiglimeccen{23}$ & $\hatcurRVeccentwosiglimeccen{24}$ \\ ~~~$\omega$ (deg) \dotfill & $\hatcurRVomegaeccen{22}$ & $\cdots$ & $\cdots$ \\ ~~~$\sqrt{e}\cos\omega$ \dotfill & $\hatcurRVrkeccen{22}$ & $\cdots$ & $\cdots$ \\ ~~~$\sqrt{e}\sin\omega$ \dotfill & $\hatcurRVrheccen{22}$ & $\cdots$ & $\cdots$ \\ ~~~$e\cos\omega$ \dotfill & $\hatcurRVkeccen{22}$ & $\cdots$ & $\cdots$ \\ ~~~$e\sin\omega$ \dotfill & $\hatcurRVheccen{22}$ & $\cdots$ & $\cdots$ \\ ~~~radial velocity jitter FEROS (\ms) \tablenotemark{g} \dotfill & \hatcurRVjitterAeccen{22} & \hatcurRVjitter{23} & \hatcurRVjitterA{24} \\ ~~~radial velocity jitter HARPS (\ms) \dotfill & \hatcurRVjitterCeccen{22} & $\cdots$ & \hatcurRVjitterC{24} \\ ~~~radial velocity jitter CORALIE (\ms) \dotfill & \hatcurRVjitterCeccen{22} & $\cdots$ & \hatcurRVjitterB{24} \\ ~~~radial velocity jitter CYCLOPS2+UCLES (\ms) \dotfill & $\cdots$ & $\cdots$ & \hatcurRVjitterD{24} \\ \sidehead{Planetary parameters} ~~~$\mpl$ ($\mjup$) \dotfill & $\hatcurPPmlongeccen{22}$ & $\hatcurPPmlong{23}$ & $\hatcurPPmlong{24}$ \\ ~~~$\rpl$ ($\rjup$) \dotfill & $\hatcurPPrlongeccen{22}$ & $\hatcurPPrlong{23}$ & $\hatcurPPrlong{24}$ \\ ~~~$\rpl$ lower limit ($\rjup$) \tablenotemark{c} \dotfill & $\cdots$ & $\hatcurPPrtwosiglowerlim{23}$ & $\cdots$ \\ ~~~$C(\mpl,\rpl)$ \tablenotemark{h} \dotfill & $\hatcurPPmrcorreccen{22}$ & $\hatcurPPmrcorr{23}$ & $\hatcurPPmrcorr{24}$ \\ ~~~$\rhopl$ (\gcmc) \dotfill & $\hatcurPPrhoeccen{22}$ & $\hatcurPPrho{23}$ & $\hatcurPPrho{24}$ \\ ~~~$\log g_p$ (cgs) \dotfill & $\hatcurPPloggeccen{22}$ & $\hatcurPPlogg{23}$ & $\hatcurPPlogg{24}$ \\ ~~~$a$ (AU) \dotfill & $\hatcurPPareleccen{22}$ & $\hatcurPParel{23}$ & $\hatcurPParel{24}$ \\ ~~~$T_{\rm eq}$ (K) \dotfill & $\hatcurPPteffeccen{22}$ & $\hatcurPPteff{23}$ & $\hatcurPPteff{24}$ \\ ~~~$\Theta$ \tablenotemark{i} \dotfill & $\hatcurPPthetaeccen{22}$ & $\hatcurPPtheta{23}$ & $\hatcurPPtheta{24}$ \\ ~~~$\log_{10}\langle F \rangle$ (cgs) \tablenotemark{j} \dotfill & $\hatcurPPfluxavglogeccen{22}$ & $\hatcurPPfluxavglog{23}$ & $\hatcurPPfluxavglog{24}$ \\ [-1.5ex] \enddata \tablenotetext{a}{ Times are in Barycentric Julian Date calculated directly from UTC {\em without} correction for leap seconds. \ensuremath{T_c}: Reference epoch of mid transit that minimizes the correlation with the orbital period. \ensuremath{T_{14}}: total transit duration, time between first to last contact; \ensuremath{T_{12}=T_{34}}: ingress/egress time, time between first and second, or third and fourth contact. } \tablecomments{ For each system we adopt the class of model which has the highest Bayesian evidence from among those tested. For \hatcurb{23} and \hatcurb{24} the adopted parameters come from a fit in which the orbit is assumed to be circular. For \hatcurb{22} the eccentricity is allowed to vary. } \tablenotetext{b}{ Reciprocal of the half duration of the transit used as a jump parameter in our Markov chain Monte Carlo (MCMC) analysis in place of $\arstar$. It is related to $\arstar$ by the expression $\zrstar = \arstar(2\pi(1+e\sin\omega))/(P\sqrt{1-b^2}\sqrt{1-e^2})$ \citep{bakos:2010:hat11}. } \tablenotetext{c}{ The grazing transits of \hatcurb{23} mean that we cannot place a strong upper limit on the impact parameter. For this system we also provide 95\% confidence lower limits on $b^{2}$, $b$ and $\rpl$, and the 95\% confidence upper limit on $i$. } \tablenotetext{d}{ Scaling factor applied to the model transit that is fit to the HATSouth light curves. This factor accounts for dilution of the transit due to blending from neighboring stars and over-filtering of the light curve. These factors are varied in the fit. } \tablenotetext{e}{ Values for a quadratic law, adopted from the tabulations by \cite{claret:2004} according to the spectroscopic (ZASPE) parameters listed in \reftabl{stellar}. } \tablenotetext{f}{ For fixed circular orbit models we list the 95\% confidence upper limit on the eccentricity determined when $\sqrt{e}\cos\omega$ and $\sqrt{e}\sin\omega$ are allowed to vary in the fit. } \tablenotetext{g}{ Term added in quadrature to the formal radial velocity uncertainties for each instrument. This is treated as a free parameter in the fitting routine. In cases where the jitter is consistent with zero, we list its 95\% confidence upper limit. } \tablenotetext{h}{ Correlation coefficient between the planetary mass \mpl\ and radius \rpl\ estimated from the posterior parameter distribution. } \tablenotetext{i}{ The Safronov number is given by $\Theta = \frac{1}{2}(V_{\rm esc}/V_{\rm orb})^2 = (a/\rpl)(\mpl / \mstar )$ \citep[see][]{hansen:2007}. } \tablenotetext{j}{ Incoming flux per unit surface area, averaged over the orbit assuming a circular geometry. } \ifthenelse{\boolean{emulateapj}}{ \end{deluxetable*} }{ \end{deluxetable} }
16
7
1607.00688
We report the discovery of three moderately high-mass transiting hot Jupiters from the HATSouth survey: HATS-22b, HATS-23b and HATS-24b. These planets add to the number of known planets in the ∼2M<SUB>J</SUB> regime. HATS-22b is a 2.74 ± 0.11 M<SUB>J</SUB> mass and 0.953_{-0.029}^{+0.048} R_J radius planet orbiting a V = 13.455 ± 0.040 sub-solar mass (M<SUB>*</SUB> = 0.759 ± 0.019 M<SUB>⊙</SUB>; R<SUB>*</SUB> = 0.759 ± 0.019 R<SUB>⊙</SUB>) K-dwarf host star on an eccentric (e = 0.079 ± 0.026) orbit. This planet's high planet-to-stellar mass ratio is further evidence that migration mechanisms for hot Jupiters may rely on exciting orbital eccentricities that bring the planets closer to their parent stars followed by tidal circularization. HATS-23b is a 1.478 ± 0.080 M<SUB>J</SUB> mass and 1.69 ± 0.24 R<SUB>J</SUB> radius planet on a grazing orbit around a V = 13.901 ± 0.010 G-dwarf with properties very similar to those of the Sun (M<SUB>*</SUB> = 1.115 ± 0.054; R<SUB>*</SUB> = 1.145 ± 0.070). HATS-24b orbits a moderately bright V = 12.830 ± 0.010 F-dwarf star (M<SUB>*</SUB> = 1.218 ± 0.036 M<SUB>⊙</SUB>; R_\star = 1.194_{-0.041}^{+0.066} R_{⊙}). This planet has a mass of 2.39_{-0.12}^{+0.21} M_J and an inflated radius of 1.516_{-0.065}^{+0.085} R_J.
false
[ "known planets", "sub-solar mass", "hot Jupiters", "tidal circularization", "J</SUB", "exciting orbital eccentricities", "a V = 13.455 ±", "a V = 13.901 ±", "M<SUB>⊙</SUB", "migration mechanisms", "1.218 ± 0.036 M<SUB>⊙</SUB", "a moderately bright V = 12.830 ±", "further evidence", "K-dwarf host star", "HATS-24b", "HATS-23b", "the Sun (M<SUB>*</SUB> = 1.115 ±", "R_J", "M<SUB>*</SUB> = 0.759 ±", "properties" ]
6.443493
14.252504
-1
3150500
[ "Mingo, B.", "Watson, M. G.", "Rosen, S. R.", "Hardcastle, M. J.", "Ruiz, A.", "Blain, A.", "Carrera, F. J.", "Mateos, S.", "Pineau, F. -X.", "Stewart, G. C." ]
2016MNRAS.462.2631M
[ "The MIXR sample: AGN activity versus star formation across the cross-correlation of WISE, 3XMM, and FIRST/NVSS" ]
77
[ "Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK", "Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK", "Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK", "School of Physics, Astronomy &amp; Mathematics, University of Hertfordshire, College Lane, Hatfield AL10 9AB, UK", "Instituto de Física de Cantabria (CSIC-UC), Avenida de los Castros, E-39005 Santander, Spain", "Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK", "Instituto de Física de Cantabria (CSIC-UC), Avenida de los Castros, E-39005 Santander, Spain", "Instituto de Física de Cantabria (CSIC-UC), Avenida de los Castros, E-39005 Santander, Spain", "CNRS, Universitè de Strasbourg, Observatoire Astronomique, 11 rue de l'Universitè, F-67000 Strasbourg, France", "Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK" ]
[ "2017A&A...597A..89P", "2017A&A...602A..79L", "2017A&A...604A..43M", "2017ApJ...837..110V", "2017ApJ...845..134W", "2017ApJ...846...64M", "2017ApJ...849...50N", "2017MNRAS.468..196T", "2017MNRAS.468..378S", "2017MNRAS.470..314H", "2017MNRAS.470.2762M", "2017MNRAS.472.2280F", "2018A&A...609A...1B", "2018A&A...618A..52R", "2018MNRAS.476.3478B", "2018MNRAS.476.5535T", "2018MNRAS.480..358W", "2018MNRAS.480..707P", "2018PhRvD..98j3007A", "2019A&A...622A..11G", "2019A&A...622A..12H", "2019A&A...622A..17S", "2019A&A...632A.109N", "2019AN....340..267L", "2019ApJ...870..104S", "2019ApJS..243...15T", "2019ApJS..245...25J", "2019MNRAS.486.4962W", "2019MNRAS.488.2701M", "2019NatAs...3..272R", "2019OAP....32...42T", "2020A&A...638A..34J", "2020A&A...640A..68I", "2020A&A...641A.136W", "2020A&A...642A.150L", "2020A&A...642A.153D", "2020ApJS..247...53K", "2020ApJS..250....7J", "2020AstSR..16....1Z", "2020FrASS...7....8L", "2020MNRAS.494.5161C", "2020MNRAS.496.1706S", "2020OAP....33...22T", "2021A&A...655A.109B", "2021ApJ...910...64K", "2021ApJ...918...65S", "2021Galax...9..102O", "2021Galax...9..122J", "2021KFNT...37c..68Z", "2021KPCB...37..149Z", "2021MNRAS.507..919L", "2022A&ARv..30....6M", "2022ApJ...931..154R", "2022ApJ...941..136S", "2022LRR....25....6M", "2022MNRAS.511.3250M", "2022MNRAS.512.3284S", "2022MNRAS.512.3858L", "2022MNRAS.514.3626C", "2022MNRAS.516.2824C", "2023A&A...671A.152M", "2023A&A...679A.101C", "2023ApJ...948L...9G", "2023ApJ...955...56K", "2023ApJS..267...37G", "2023JApA...44...13D", "2023MNRAS.518.4290S", "2023MNRAS.524.3958E", "2024A&A...685A.108E", "2024ApJ...965...17D", "2024ApJ...966...53F", "2024EPJC...84..444L", "2024MNRAS.527.3006G", "2024MNRAS.527.3436P", "2024MNRAS.529.1472C", "2024PASA...41...27G", "2024arXiv240513650C" ]
[ "astronomy" ]
8
[ "galaxies: active", "galaxies: starburst", "infrared: galaxies", "radio continuum: galaxies", "X-rays: galaxies", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stw1826", "10.48550/arXiv.1607.06471" ]
1607
1607.06471_arXiv.txt
Over the last few decades, and in particular in the last 10 to 15 years, our understanding of active galactic nuclei (AGN), their underlying physical mechanisms, their environments, and their observational properties, has greatly increased. Although the unification model proposed by \citet{Antonucci1993} still holds true in many aspects, subsequent revisions \citep[see e.g.][]{Netzer2015} illustrate what we have learned about the structure of the obscuring torus, the mechanisms that provide feedback, the variability timescales involved, and where the radio-loud sources fit (or do not fit) in the grand AGN unification scheme. We are living in what could be considered a golden era of surveys, which allow us, for the first time, to construct large, consistent, multiwavelength samples of AGN with the potential to push our understanding of these objects even further. Although only $\sim$10--20 per cent of the AGN we observe are classified as radio-loud, recent evidence shows that jets and lobes could be far more ubiquitous than we previously thought. There is an increasingly large number of Seyfert galaxies, and even QSOs, where jets and lobes, or excess radio emission, have been detected \citep[e.g.][]{Hota2006,Gallimore2006,DelMoro2013,Singh2015,Harrison2015}, throwing into question the radio-loud/quiet classification, which, being based on optical (B band) to radio (5 GHz) flux ratios \citep[e.g.][]{Kellerman1989}, classifies most of these objects as radio-quiet. This `jet mode' or `radio mode' is fundamental to our understanding of the AGN/host relationship, not only for very powerful sources in clusters, where the jet-driven shocks can offset radiative cooling of the gas \citep[e.g.][]{McNamara2012,HLarrondo2015}, but especially for low power sources ($L_{1.4 GHz} \le 10^{23}$ W Hz$^{-1}$ sr$^{-1}$), because it is in these systems that the effect of the AGN on the surrounding interstellar gas (on 10--100 kpc scales) can have the largest potential impact on the evolution and star formation history of the host galaxy \citep[e.g.][]{Cattaneo2009,Croston2011,Mingo2011,Mingo2012}. Radio-loud sources are also useful in that they allow us to unequivocally identify sources in the radiatively inefficient accretion regime \citep{Narayan1995}. This population, originally identified as low excitation radio galaxies (LERGs) by \citet{Hine1979}, lacks the `traditional' AGN disc and torus, shows very low Eddington rates \citep{Hardcastle2007b,Hardcastle2009,DeGasperin2011,Best2012,Mingo2014,Paggi2016}, and seems to be channeling most of the gravitational energy into jets rather than radiative output, in a similar manner to the low/hard state of low mass X-ray binaries \citep[see e.g. the review by][]{Fender2014}. In the optical, radiatively inefficient sources are typically classified as LINERs (low ionisation nuclear emission line regions) as initially proposed by \citet{Heckman1980}, or even appear as fully `quiescent' (i.e. not containing an AGN) galaxies \citep[see e.g.][]{Kimball2008}. This classification, however, is misleading, in the sense that other processes such as shocks or emission from an old stellar population can also produce low ionisation spectra \citep[see e.g.][and references therein]{Balmaverde2015}. Therefore, finding low ionisation optical emission lines does not guarantee the presence of a radiatively inefficient AGN, while finding active radio jets in an otherwise `quiescent' looking galaxy does. As radiatively inefficient AGN only produce soft X-rays related to the jet \citep[e.g.][]{Hardcastle1999}, their typical X-ray luminosity is $10^{39}-10^{41}$ erg/s, which precludes them from being included in most X-ray selected AGN surveys. Recent results show that the interplay between AGN activity, outflows, and star formation may be more complex than we previously thought, and fundamental to understanding galaxy evolution and black hole growth \citep[e.g.][]{Alexander2012,Magliocchetti2014,Davies2014}. Although we are beginning to better understand the transition between the regimes in which AGN and star formation activity dominate, and how radio AGN activity, in particular, affects star formation \citep[e.g.][]{Smolcic2009,Dicken2012,DelMoro2013,Hardcastle2013b,Kalfountzou2014,Villarroel2014,Gurkan2015,Rawlings2015,Hardcastle2016,Drouart2016,Tadhunter2016}, there is still a distinct lack of agreement on how and when AGN activity influences star formation \citep{Harrison2012,Ishibashi2012,Symeonidis2013,Symeonidis2014,AlonsoHerrero2013,Heckman2014,Balmaverde2016,Brusa2015,Rosario2013,Rosario2015,Stanley2015,Bernhard2016,Alberts2016}. Although the large timescales involved probably cause part of this confusion \citep{Georgakakis2008,Wild2010,RamosAlmeida2013,Best2014}, and it is clear that we still do not fully understand long AGN variability timescales \citep[see e.g.][]{Hickox2014}, it is also true that dedicated samples that encompass sources in both regimes, as well as the transition, still tend to be limited either in wavelength, scope, redshift, or size. Obtaining large multiwavelength samples of radio-loud AGN is challenging for several reasons, the main two being the extended nature of radio emission and the low sky density of radio-loud AGN, and the number of sources decreases rapidly if selections in more than two bands are required. These surveys also tend to focus on particular populations of radio-loud AGN (or star-forming galaxies). There is a wealth of on-going and upcoming instruments and surveys that will open a wide field of potential exploration in both fields: LOFAR, SKA, e-MERLIN, JVLA, in the radio; e-Rosita, and Athena in the X-rays, CTA in the gamma-ray band, LSST, and JWST and ALMA at infrared and sub-mm wavelengths, respectively. Now is the perfect time to assess which questions our current data can and cannot answer, to set a framework and potential diagnostic tools for the next generation of results. The ARCHES FP7 collaboration\footnote{\url{http://www.arches-fp7.eu/}} is a project dedicated to fully exploiting the capabilities of the 3XMM catalogue of X-ray sources, by creating multiwavelength products (cross-correlated catalogues and tools, spectral energy distributions, and a cluster catalogue and finder tool). As part of this collaboration, we have built and describe in this paper the MIXR sample: a systematic, large sample of sources detected in the Mid-IR (WISE all-sky survey), X-rays (3XMM DR5) and Radio (FIRST/NVSS). By requiring a detection in all three bands, we find a wide range of populations: from radiatively inefficient (LERG/LINER) systems in otherwise quiescent galaxies, to low luminosity Seyfert-like sources where the host emission dominates in some bands, to nearby starburst objects, to high luminosity radio-loud and radio-quiet Seyferts and QSOs. The MIXR sample allows us to derive efficient diagnostics for star formation and AGN activity (both radiatively efficient, as seen in `traditional' AGN, and radiatively inefficient, as seen in LERG/LINER), even in host-dominated sources that are normally considered quiescent and discarded from most mid-IR and X-ray AGN samples. We also test the radiative (luminosity) versus kinetic (jet) output in our AGN, to explore the extent and possible causes for the scatter we observed in \citet{Mingo2014}, in contradiction with the well-known correlation of \citet{Rawlings1991}. Our analysis also helps us pinpoint several sources of bias that affect selections performed in one or more of the bands we use, helping us better understand what AGN populations are included and excluded in each selection. In section \ref{Data} we discuss in detail the MIXR sample construction. In section \ref{Diagnostics} we use WISE colours to pre-classify the sources, and carry out a series of early diagnostics to test these classifications, using hardness ratios, radio versus X-ray `loudness', and flux/magnitude diagrams. In section \ref{z} we add redshift information from SDSS, which we use in section \ref{LDiag} to derive luminosities for the MIXR sources, and extend our diagnostics to verify the underlying type of activity for the MIXR sources. In section \ref{RL_RQ_section} we re-classify the sources based on their activity (radio-quiet and radio-loud AGN, including LERGs/LINERs, and galaxies). For sections \ref{Eddington} and \ref{Power} we focus on the AGN, assessing their Eddington rates and their radiative versus kinetic (jet) output, to highlight the strengths and limitations of current surveys, and address some of the open questions on AGN variability and its impact on the AGN/host relationship. For this work we have used the latest cosmological values released by the Planck collaboration \citep{Planck2015}: $H_{0} = 67.74$ km/s/Mpc, $\Omega_{m} = 0.3089$, and $\Omega_{\Lambda} = 0.6911$. The catalogue we describe in this paper is available on-line for download at \url{http://www.arches-fp7.eu/index.php/tools-data/downloads/mixr-catalogue} and will be made available on VizieR.
We have used the ARCHES xmatch statistical tool to create a large cross-correlated sample of AGN and star-forming galaxies, using the largest, most uniform catalogues available in the Mid-IR (WISE), X-rays (3XMM) and Radio (FIRST+NVSS) bands. The MIXR sample we thus obtain provides efficient and broad-reaching diagnostic tools to classify sources based on their type of activity (radio-loud and radio-quiet AGN, and star formation), even in the absence of redshifts. The techniques we have developed for MIXR can be used to triage sources for any extragalactic sample with measurements that can be translated to these bands, paving the way for classification techniques that will allow us to fully exploit the vast amounts of data that the next generation of instruments will make available. We pre-classify our sources based on their mid-IR colours, using the WISE colour/colour plot and the results of \citet{Lake2012}, as elliptical, spiral, starburst and AGN sources. While these initial classifications provide a general idea of the type of underlying activity we can expect in our sources, there is a great deal of overlap between populations (see Table \ref{Activity}). We use first flux and magnitude plots, and then luminosity plots, to triage our sources based on their emission in each band, clearly separating star-forming, non-active galaxies \citep[for which we recover the radio/IR correlation of][]{DeJong1985,Appleton2004,Garrett2015} from radio-loud AGN, both of the radiatively efficient and inefficient \citep[LERG/LINER, see][]{Narayan1995} varieties, and from radio-quiet AGN, where the bulk of the radio emission we detect is produced by star formation, or particle acceleration in shocks \citep{Zakamska2016,Nims2015}, but which could also host minor jets and lobes. Our results show that WISE-colour selected AGN samples are heavily biased against Seyfert-type, moderate- to low-luminosity AGN. This selection bias occurs in two ways: WISE colour cuts such as those of \citet{Assef2010,Mateos2012,Stern2012} are very efficient at selecting clean samples of luminous AGN, but necessarily omit those sources where the host contributes a substantial fraction of the total emission, as only with additional proxies, such as radio emission, it is possible to distinguish between non-active galaxies and AGN at low luminosities; the WISE W3 (and W4) band is also not sensitive enough to detect faint AGN at redshifts beyond $\sim0.1$, which is particularly detrimental to radio-loud Seyfert-like sources, as these tend to appear at higher redshifts than radio-quiet sources of similar bolometric luminosity. In fact, our sample size is cut by $\sim40$ per cent simply by imposing requirement for a signal/noise of 3 in W3, with radio-loud AGN suffering the bulk of the cut (we lose another $40$ per cent of the sample when requiring SDSS redshifts for the second part of our diagnostics). We find that RL and RQ AGN of similar bolometric luminosities and Eddington rates are found at different redshifts, with the RL sources being found at slightly larger $z$, and our sources become `radio-louder' with increasing redshift, up to our detection limit. Our sample is biased against RQ AGN, as we require a radio detection, limiting the redshift more quickly for RQ sources than for RL ones. As a consequence, when considering both populations as a whole, our RL AGN are more luminous and have larger Eddington rates. This is clearly very likely to be a selection effect, and it illustrates one of the easiest causes of bias that can be incurred when comparing RL and RQ AGN: it is not enough to match both samples exclusively on luminosity, redshift, or Eddington rate; all these variables (plus their environments) must be taken into account to ensure that we are comparing like with like. Perhaps the most crucial result of this work is the confirmation of the scatter we observed in the 2Jy and 3CRR sources \citep{Mingo2014} between their radiative (bolometric luminosity) and kinetic (jet) output, in contradiction with the tightness of the long-standing correlation of \citet{Rawlings1991}. These two quantities must have a common underlying mechanism, as they are both tied to accretion, but either jet regulating mechanisms \citep[e.g.][]{Done2014,Tchekhovskoy2011,Hawley2015}, dispersion in the jet power/L$_{radio}$ relationship, long-term AGN variability \citep[e.g.][]{Hickox2014,Stanley2015}, or a combination of all three, are introducing the 4--5 order of magnitude scatter we observe in our plots. Given what we know about the potential impact of small-scale radio sources on the energetics of their hosts \citep[e.g.][]{Croston2007}, and the recently found coexistence of radiative winds and radio outflows \citep{Nesvadba2008,Harrison2015}, which has no parallel in X-ray binaries, we may need to reassess what we know about the interplay between AGN activity and star formation. If both the bolometric luminosity and the jet output of an AGN can vary by 3--5 orders of magnitude in the space of a few Myr, and we cannot extrapolate their life cycles from those we know about from LMXRB, how can we analyse the interplay between AGN activity and star formation? Although star formation occurs on longer timescales, jet-driven shocks can carry enough energy, far enough into the ISM, to potentially change the course of an on-going episode of star formation. However, we may not be able to detect whether the star formation rates we measure are influenced by consecutive radio outflows that have long faded out, and are completely unrelated to the current radiative and jet properties of the AGN because of the short timescale and wide range of AGN variability. Radio-loud Seyferts may hold the key both to understanding the details of the jet-ISM interaction, and the mechanisms regulating the jet. Some of these sources can produce jet outputs similar to those of luminous QSOs, but at values of $L_{bol}$ and $L_{Edd}$ that are orders of magnitude lower. Unfortunately, these are exactly the sources that W3 is not sensitive enough to reliably detect, as they have similar mid-IR luminosities to those of radio-quiet Seyferts at $z\sim$0.1--0.3, but are far more distant ($z\sim$0.8--1). We clearly need to better understand the life cycles of the radio-loud phase of AGN, both from a theoretical and from an observational point of view. A sensible first step might be to assess the fraction and properties of sources with radio emission in samples that do not use radio selections, supplemented with dedicated studies of moderate- to low-luminosity AGN, to establish larger samples that we can systematically study from a broad perspective that includes the hosts. We could also focus on samples such as those of e.g. \citet{Lonsdale2015}, for which the radio emission is compact, and the AGN and star formation are acting on similar spatial and time scales. We have also seen that it might be time to revisit and redefine the radio-loud/quiet classifications, as we have shown that the distribution of sources in terms of L$_{1.4 GHz}$/L$_{2-10 keV}$ displays a gradual transition between both regimes, rather than a dichotomy, showing that in many AGN there is a coexistence of two complicatedly interwoven regimes (kinetic and radiative), both with the potential to influence the host galaxy in different ways. \subsection*{Acknowledgements} We thank the anonymous referee for insightful and constructive comments that greatly improved the paper. We also thank the TOPCAT developer Mark Taylor for programming and releasing a software patch so that we could improve our Figures as the referee requested, and Dr B. Punsly, for allowing us to plot the data of \citet{Punsly2011} again in Figs. \ref{Punsly} to \ref{Punsly2_rej}. This work has made use of data/facilities and financial support from the ARCHES project (7th Framework of the European Union n$^{\circ}$ 313146). SM, FJC and AR acknowledge financial support by the Spanish Ministry of Economy and Competitiveness through grant AYA2012-31447, which is partly funded by the FEDER programme. FJC also acknowledges financial support through grant AYA2015-64346-C2-1-P (MINECO/FEDER). This work is based on observations obtained with \textit{XMM-Newton}, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. It also makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration. We acknowledge the use of the FIRST and NVSS catalogues, provided by the NRAO. Optical magnitudes and redshifts were obtained from the Sloan Digital Sky Survey Data Release 12. Funding for the Sloan Digital Sky Survey IV has been provided by the Alfred P. Sloan Foundation and the Participating Institutions. SDSS-IV acknowledges support and resources from the Center for High-Performance Computing at the University of Utah. The SDSS web site is \url{www.sdss.org}.
16
7
1607.06471
We cross-correlate the largest available mid-infrared (Wide-field Infrared Survey Explorer - WISE), X-ray (3XMM) and radio (Faint Images of the Radio Sky at Twenty centimetres+NRAO VLA Sky Survey) catalogues to define the MIXR sample of AGN and star-forming galaxies. We pre-classify the sources based on their positions on the WISE colour/colour plot, showing that the MIXR triple selection is extremely effective to diagnose the star formation and AGN activity of individual populations, even on a flux/magnitude basis, extending the diagnostics to objects with luminosities and redshifts from SDSS DR12. We recover the radio/mid-IR star formation correlation with great accuracy, and use it to classify our sources, based on their activity, as radio-loud and radio-quiet active galactic nuclei (AGN), low excitation radio galaxies/low ionization nuclear emission line regions, and non-AGN galaxies. These diagnostics can prove extremely useful for large AGN and galaxy samples, and help develop ways to efficiently triage sources when data from the next generation of instruments becomes available. We study bias in detail, and show that while the widely used WISE colour selections for AGN are very successful at cleanly selecting samples of luminous AGN, they miss or misclassify a substantial fraction of AGN at lower luminosities and/or higher redshifts. MIXR also allows us to test the relation between radiative and kinetic (jet) power in radio-loud AGN, for which a tight correlation is expected due to a mutual dependence on accretion. Our results highlight that long-term AGN variability, jet regulation, and other factors affecting the Q/L<SUB>bol</SUB> relation, are introducing a vast amount of scatter in this relation, with dramatic potential consequences on our current understanding of AGN feedback and its effect on star formation.
false
[ "low excitation radio galaxies", "non-AGN galaxies", "AGN feedback", "luminous AGN", "AGN", "star formation", "low ionization nuclear emission line regions", "SDSS DR12", "non-AGN", "lower luminosities", "VLA Sky Survey", "radio", "higher redshifts", "radio-loud AGN", "AGN activity", "samples", "large AGN and galaxy samples", "WISE", "sources", "the radio/mid-IR star formation correlation" ]
14.402646
7.291813
119
12472558
[ "Pović, Mirjana", "Márquez, Isabel", "Netzer, Hagai", "Masegosa, Josefa", "Nordon, Raanan", "Pérez, Enrique", "Schoenell, William" ]
2016MNRAS.462.2878P
[ "Star formation and AGN activity in the most luminous LINERs in the local universe" ]
16
[ "Instituto de Astrofísica de Andalucía (IAA-CSIC), E-18008 Granada, Spain", "Instituto de Astrofísica de Andalucía (IAA-CSIC), E-18008 Granada, Spain", "School of Physics and Astronomy and the Wise Observatory, The Raymond and Beverly Sackler Faculty of Exact Sciences, Tel-Aviv University, Tel-Aviv 69978, Israel", "Instituto de Astrofísica de Andalucía (IAA-CSIC), E-18008 Granada, Spain", "School of Physics and Astronomy and the Wise Observatory, The Raymond and Beverly Sackler Faculty of Exact Sciences, Tel-Aviv University, Tel-Aviv 69978, Israel", "Instituto de Astrofísica de Andalucía (IAA-CSIC), E-18008 Granada, Spain", "Instituto de Astrofísica de Andalucía (IAA-CSIC), E-18008 Granada, Spain" ]
[ "2017A&A...603A.131H", "2017FrASS...4...34M", "2017MNRAS.470.1687B", "2017MNRAS.471.3226M", "2018MNRAS.480.1106C", "2018MNRAS.480.3993B", "2019ApJ...871..190M", "2019MNRAS.485..452M", "2019arXiv190102158M", "2020MNRAS.498.4345Z", "2021IAUS..356....3P", "2021IAUS..356..147M", "2021IAUS..356..323M", "2021MNRAS.505.4289P", "2022MNRAS.513.4494M", "2024MNRAS.531..199O" ]
[ "astronomy" ]
5
[ "galaxies: active", "galaxies: nuclei", "galaxies: star formation", "Astrophysics - Astrophysics of Galaxies" ]
[ "1980A&A....87..152H", "1981PASP...93....5B", "1983ApJ...264..105F", "1985MNRAS.213..841T", "1989ApJ...345..245C", "1992ApJ...388..310K", "1994ApJ...429..582C", "1996ARA&A..34..749S", "1997ApJ...490..202D", "1997ApJS..111..377W", "1997ApJS..112..315H", "1998ApJ...500..525S", "1999ApJ...527...54B", "2000ApJ...537L..85R", "2001ApJ...556..121K", "2001ApJ...556..562C", "2002ApJ...574..740T", "2003ApJ...583..159H", "2003ApJ...584...76C", "2003MNRAS.341...33K", "2003MNRAS.341...54K", "2003MNRAS.344.1000B", "2003MNRAS.346.1055K", "2004AJ....127.2002K", "2004ApJ...605..105C", "2004MNRAS.351.1151B", "2005A&A...435..521N", "2005ASPC..347...29T", "2005MNRAS.358..363C", "2005MNRAS.362.1143M", "2006A&A...460...45G", "2006ApJ...642..775M", "2006ApJ...643...14S", "2006ApJ...648..281Y", "2006ApJS..162...38A", "2006MNRAS.370..721M", "2006MNRAS.371..972S", "2006MNRAS.372..961K", "2007A&A...468...33E", "2007ApJ...660L..43N", "2007ApJ...666..870C", "2007ApJ...670..156D", "2007ApJS..172..150S", "2008ARA&A..46..475H", "2008ApJ...682..319B", "2008MNRAS.391L..29S", "2009A&A...506.1107G", "2009ApJ...704.1570G", "2009MNRAS.399.1907N", "2010A&A...518L..25R", "2010A&A...519A..40A", "2010ApJ...712.1287L", "2010ApJ...713..115G", "2010MNRAS.403.1036C", "2011A&A...529A.126C", "2011ApJ...738..106W", "2011MNRAS.413.1687C", "2012A&A...540A.109S", "2012A&A...545A..45R", "2012A&A...546A..58R", "2012ApJ...744..159L", "2012ApJ...747...61Y", "2012ApJ...753..155T", "2012ApJ...754L..29W", "2012ApJ...755..131R", "2012Natur.485..213P", "2013A&A...557A..86C", "2013A&A...558A..43S", "2013ApJ...764L...1P", "2013ApJ...770...57B", "2013ApJ...775...41A", "2013ApJ...778...23G", "2013peag.book.....N", "2014ARA&A..52..589H", "2014ApJ...782....9H", "2014ApJ...795..104W", "2015ApJ...801...87B", "2015ApJ...806..187A", "2015MNRAS.452.1841S", "2015MNRAS.453..591S", "2016A&A...590A..44G", "2016ApJ...819..123N", "2016MNRAS.455L..82L" ]
[ "10.1093/mnras/stw1842", "10.48550/arXiv.1607.03915" ]
1607
1607.03915_arXiv.txt
\label{sec_intro} \indent \indent Low Ionization Nuclear Emission line Regions (LINERs) are the most common active galactic nuclei (AGN), with numbers that exceed those of 'high ionization AGN' (type-I and type-II Seyfert galaxies and quasars) \citep{heckman80,ho08,heckman14}. At least in the local universe they make up 1/3 of all galaxies and 2/3 of AGN population \citep{kauffmann03a,yan06,ho08}. LINERs are normally classified by their narrow emission line ratios, e.g. [OIII]$\lambda$5007/H$\beta$, [NII]$\lambda$6584/H$\alpha$, and [OI]$\lambda$6300/H$\alpha$ \citep{baldwin81,kauffmann03a,stasinska06,kewley06}. In general, they have lower luminosities than Seyfert galaxies, but there is a big overlap between the groups in terms of properties like stellar mass, X-ray and radio luminosity, etc. \citep{ho08,netzer09,leslie16}.\\ \indent Different mechanisms were proposed to explain the nature of LINERs. This includes shock excitation \citep[e.g.][]{dopita97,nagar05}, photoionisation by young, hot, massive stars \citep{terlevich85}, photoionisation by evolved post-asymptotic giant branch (pAGB) stars \citep[e.g.][]{stasinska08,annibali10,cid11,yan12,singh13}, and photoionisation by a central low-luminosity AGN \citep[e.g.][]{ferland83,ho08,gonzalezmartin06}. The first two proposals failed to explain the properties of large samples of LINERs. The third possibility of pAGB stars was suggested for LINERs with the weakest emission lines, located in galaxies with predominately old stars. They can be distinguished from strong-line LINERs using the equivalent widths (EW) of their emission lines, e.g., EW([OIII]$\lambda$5007)\,$<$\,1\,\AA\,\citep{capetti11} or EW(H$\alpha$)\,$<$\,3\,\AA\,\citep{cid11}. Several works however questioned this possibility, arguing that a population that is less luminous and more numerous than pAGB stars would be needed to produce the luminosities observed in weak LINERs \citep{brown08,rosenfield12,heckman14}. However, most LINERs are powered by an AGN, especially those with stronger emission lines (e.g., EW(H$\alpha$)\,$>$\,3\AA) and unresolved hard X-ray emission \citep[e.g., ][and references therein]{gonzalezmartin06,gonzalezmartin09a,gonzalezmartin09b,heckman14}. Like other AGNs, LINERs can be divided into type-I (broad and narrow emission lines) and type-II (only narrow emission lines). Their emission lines are characterised by lower levels of ionization than in Seyferts, and their normalized accretion rates (Eddington ratio) are 1-5 orders of magnitude smaller.\\ \indent The best studied nearby LINERs \citep[e.g.][]{ho97,ho08,kauffmann03a,leslie16} are found in nuclei of galaxies with little or no evidence of active star formation (SF). They are usually characterised as being hosted by massive early-type galaxies (rarely spirals), and massive black holes in their centres, old stellar populations, small amounts of gas and dust, with low extinctions. Such LINERs show weak and small-scale radio jets \citep{ho08,heckman14}. \\ \indent \cite{tommasin12} studied SF in LINERs from the COSMOS field at z\,$\sim$\,0.3 using Herschel/PACS observations. They showed that: a) The SF luminosities of 34 out of 97 high luminosity LINERs are on average 2 orders of magnitude higher than SF luminosities of lower AGN luminosity, nearby LINERs. b) Even if assumed that all the observed H$\alpha$ flux is due to SF (a wrong assumption since much of it must be due to AGN excitation) it is still impossible to recover the SF rate (SFR) indicated by the FIR observations. Given this result, we suspect that active SF in LINER host galaxies has escaped the attention of most earlier studies that focused on the innermost part of nearby galaxies. In this work we focus on the most luminous LINERs in the local (0.04\,$<$\,z\,$<$\,0.11) universe and study their SF and AGN activity, in order to understand the LINER phenomenon in relation to star-forming galaxies and to compare their properties with those of the LINERs at z\,$\sim$\,0.3. Many properties of these sources are known from SDSS spectroscopy and/or GALEX observations, e.g., emission line luminosities, locations on the BPT diagrams, SFRs based on Dn4000 estimations, etc. Unfortunately, the 3\,arc-sec SDSS fibre does not allow to resolve the nuclear region and hence to separate AGN excited from SF excited emission lines. The goals of the present study are to carry out a detailed, ground based spectroscopy of the central regions of the most luminous LINERs, and to measure, together with Herschel and IRAS FIR data, their SFRs in a careful way. \\ \indent The paper is organised as follows: in Section~\ref{sec_sample_selection} we describe the sample selection. Reduction procedure for our new spectroscopic data, together with our own or archival FIR data are described in Section~\ref{sec_observations}. In Section~\ref{sec_analysis} we summarise all our measurements, including spectral fittings, emission line and extinction measurements, and estimations of Dn4000 and H$\delta$ indices, AGN luminosities, and SFRs. The main results are presented in Section~\ref{sec_results_discussion} where we discuss the general properties of the most luminous LINERs in the local universe, co-evolution between the SF and AGN activity, and the location of our sample on the main sequence of SF galaxies. \\ \indent We assumed the following cosmological parameters throughout the paper: $\Omega_{\Lambda}$\,=\,0.7, $\Omega_{M}$\,=\,0.3, and H$_0$\,=\,70 km s$^{-1}$ Mpc$^{-1}$. \section[]{Sample selection} \label{sec_sample_selection} \indent \indent The sources were initially selected from the SDSS/DR4 \citep{kauffmann03,brinchmann04} catalogue in Garching MPA-JHU based on the Sloan Digital Sky Survey (SDSS\footnote{http://www.sdss.org/}) DR4 data \citep[][and references therein]{adelman06}. LINERs were first selected using both [NII]$\lambda$6584/H$\alpha$ and [OI]$\lambda$6300/H$\alpha$ criteria of \cite{kewley06}. Taking into account the completeness of the SDSS survey, only LINERs with 0.04\,$<$\,z\,$<$\,0.11 were selected \citep{netzer09}. To eliminate LINERs ionised by pAGB stars we selected only those galaxies with H$\alpha$ equivalent width EW(H$\alpha$)\,$>$\,2.5\AA ~\citep{cid11}. \\ % \indent The next step was the selection of the most luminous LINERs within the chosen redshift interval. We measured first their AGN luminosity (LAGN) using the [OIII]$\lambda$5007 and [OI]$\lambda$6300 method of \cite{netzer09} (see Section~\ref{sec_agn_lum}). The lines were initially corrected for reddening using the observed H$\alpha$/H$\beta$ ratio and assuming galactic extinction (see Sections~\ref{sec_emission_measure} and \ref{sec_agn_lum}). We selected a certain, statistically sufficient, fraction of 147 luminous LINERs with logLAGN\,$>$\,44.3\,ergs/sec. We call these sources 'LLINERs'. Out of these sources we selected a luminosity limited sample of 47 galaxies with SF luminosity LSF\,$>$\,43.3\,ergs/sec, where LSF is based on the Dn4000 index (see Section~\ref{sec_sfr}). Of those, we were able to obtain the optical spectra for 42 LINERs and Herschel/PACS data for 6 sources. We refer to these 42 most luminous LINERs in terms of both AGN and SF luminosity as 'MLLINERs'. All observed MLLINERs are listed in Table~\ref{tab_obs}, where we provide the basic information about their properties. \\ \indent Figure~\ref{fig_sample_selection} shows the position in the LAGN vs. LSF plane of the initially classified LINERs in the selected redshift range (black dots), and the final selected sample of MLLINERs (blue squares). Using the SDSS spectroscopy we estimated the AB continuum magnitude at 6500\,\AA\,(m6500). We used these magnitudes to divide the sample into 'faint' and 'bright' galaxies (m6500\,$>$\,17.2\,mag and m6500\,$<$\,17.2\,mag, respectively). These groups are marked with F or B in Table 1. We use this classification only for observational purposes. Figs~\ref{fig_thumbnails_part1} and \ref{fig_thumbnails_part2} show SDSS colour images of all MLLINERs. \begin{figure} \centering \includegraphics[width=0.49\textwidth,angle=0]{HAGAI_mongo_herschel_selection_for_liner_paper.ps} \caption{The entire 0.04\,$<$\,z\,$<$\,0.11 SDSS/DR4 LINER sample used in this work (small black squares) and the sub-sample used for the Herschel proposal and the follow up spectroscopy (large blue squares). The dashed lines mark the lower limits on LAGN and LSF (based on Dn4000 index) used for the selection of the targets. \label{fig_sample_selection}} \end{figure} \begin{table*} \small \begin{center} \caption{Summary of observations. \label{tab_obs}} \begin{tabular}{| c | c | c | c | c | c | c | c | c | c | c | c | c |} \hline \textbf{ID}&\textbf{RA}&\textbf{DEC}&\textbf{z}&\textbf{m6500}&\textbf{morph}&\textbf{Date}&\textbf{seeing}&\textbf{pos. ang.}&\textbf{texp\_b}&\textbf{texp\_r}& \textbf{Area$_{nuc}$}& \textbf{IR data}\\ &\textbf{[deg]}&\textbf{[deg]}&&\textbf{AB [mag]}&&&\textbf{[arc-sec]}&\textbf{[deg]}&\textbf{[sec]}&\textbf{[sec]}&\textbf{[arc-sec$^2$]}& \\ \hline F01 & 47.499332 & 0.29955 & 0.098 & 18.53 & S & 02/11/2013 & 1.2 & PA & 3\,$\times$\,3000.0 & 3\,$\times$\,3000.0 & 3.6 & \\ F02 & 115.434586 & 21.18252 & 0.098 & 17.96 & E & 06/03/2014 & 1.2 & 314 & 3\,$\times$\,3000.0 & 3\,$\times$\,3000.0 & 3.6 & 1, 2 \\ F03 & 131.35008 & 39.245438 & 0.109 & 17.63 & P & 08/03/2014 & 1.3 & 201 & 3\,$\times$\,2000.0 & 3\,$\times$\,2000.0 & 3.9 & \\ F04 & 129.59967 & 49.04478 & 0.101 & 17.58 & P & 08/03/2014 & 1.3 & 206 & 3\,$\times$\,2400.0 & 3\,$\times$\,2400.0 & 3.9 & \\ F06 & 144.995 & 34.96791 & 0.104 & 17.63 & E & 09/03/2014 & 1.4 & 220 & 3\,$\times$\,2000.0 & 3\,$\times$\,2000.0 & 4.2 & 2 \\ F07 & 138.23363 & 46.8671 & 0.051 & 17.27 & E & 09/03/2014 & 1.4 & 338 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 4.2 & \\ F09 & 170.5683 & 54.6951 & 0.105 & 17.50 & S & 03/05/2014 & 1.4 & 149 & 3\,$\times$\,2000.0 & 3\,$\times$\,2000.0 & 4.2 & 2 \\ F12 & 182.36954 & 11.030761 & 0.107 & 17.21 & S & 02/05/2014 & 1.6 & 410 & 3\,$\times$\,2400.0 & 3\,$\times$\,2400.0 & 6.0 & 2 \\ F13 & 183.83566 & 5.533633 & 0.082 & 18.09 & E & 05/05/2014 & 1.2 & 120 & 3\,$\times$\,3000.0 & 3\,$\times$\,3000.0 & 3.6 & \\ F14 & 180.15637 & 4.530397 & 0.094 & 17.54 & S & 04/05/2014 & 1.2 & 265 & 3\,$\times$\,2000.0 & 3\,$\times$\,2000.0 & 3.6 & 2 \\ F15 & 203.8548 & 45.891083 & 0.092 & 17.42 & E & 03/05/2014 & 1.4 & 239 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 4.2 & \\ F16 & 255.87796 & 20.849482 & 0.08 & 18.43 & ? & 26/07/2014 & 1.0 & 184 & 3\,$\times$\,3600.0 & 3\,$\times$\,3600.0 & 3.0 & 2 \\ F17 & 259.5603 & 64.29323 & 0.104 & 17.78 & E & 06/03/2014 & 1.2 & 213 & 3\,$\times$\,2400.0 & 3\,$\times$\,2400.0 & 3.6 & 1, 2 \\ F19 & 316.2105 & 0.358728 & 0.091 & 17.90 & ? & 25/07/2014 & 1.0 & 205 & 3\,$\times$\,2800.0 & 3\,$\times$\,2800.0 & 3.0 & 1 \\ F20 & 333.30197 & 13.3283 & 0.103 & 18.53 & P & 27/07/2014 & 1.3 & 127 & 3\,$\times$\,3600.0 & 3\,$\times$\,3600.0 & 3.9 & \\ F21 & 342.84195 & -8.956378 & 0.08 & 17.50 & E & 28/07/2014 & 1.6 & 241 & 3\,$\times$\,3000.0 & 3\,$\times$\,3000.0 & 4.8 & \\ F22 & 358.20468 & 14.04565 & 0.096 & 18.02 & ? & 29/07/2014 & 1.2 & 238 & 3\,$\times$\,3200.0 & 3\,$\times$\,3200.0 & 3.6 & \\ F23 & 9.282583 & 0.410139 & 0.081 & 17.42 & ? & 30/07/2014 & 1.4 & 260 & 3\,$\times$\,2800.0 & 3\,$\times$\,2800.0 & 4.2 & \\ F24 & 23.73075 & -8.710756 & 0.092 & 18.02 & P & 09/10/2013 & 1.2 & PA & 3\,$\times$\,3000.0 & 3\,$\times$\,3000.0 & 3.6 & 2 \\ B01 & 53.543957 & 1.103353 & 0.048 & 17.17 & S & 31/10/2013 & 1.5 & PA & 3\,$\times$\,1600.0 & 3\,$\times$\,1600.0 & 5.6 & \\ B02 & 124.66104 & 23.48597 & 0.103 & 16.90 & P & 07/03/2014 & 1.2 & 315 & 3\,$\times$\,1700.0 & 3\,$\times$\,1700.0 & 3.6 & 1 \\ B03 & 129.57721 & 33.57853 & 0.062 & 16.79 & P & 06/03/2014 & 1.2 & 274 & 3\,$\times$\,1600.0 & 3\,$\times$\,1600.0 & 3.6 & 1, 2 \\ B04 & 133.79796 & 0.219117 & 0.101 & 16.90 & E & 10/03/2014 & 1.3 & 255 & 3\,$\times$\,1700.0 & 3\,$\times$\,1700.0 & 3.9 & \\ B05 & 141.73837 & 8.630544 & 0.106 & 17.09 & S & 10/03/2014 & 1.3 & 180 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 3.9 & 2 \\ B06* & 160.26555 & 11.096189 & 0.053 & 16.50 & ? & 01/05/2013 & 0.9 & PA & 4\,$\times$\,900.0 & 3\,$\times$\,900.0 & 3.0 & \\ B07 & 165.55441 & 66.1674 & 0.078 & 17.17 & P & 06/03/2014 & 1.2 & 245 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 3.6 & 1 \\ B08* & 170.29817 & -0.293878 & 0.098 & 17.11 & E & 03/05/2013 & 0.9 & PA & 4\,$\times$\,1200.0 & 4\,$\times$\,900.0 & 3.0 & \\ B09 & 171.66946 & -1.6938 & 0.046 & 15.93 & E & 10/03/2014 & 1.3 & 290 & 3\,$\times$\,1200.0 & 3\,$\times$\,1200.0 & 3.9 & \\ B10 & 183.72675 & 1.916183 & 0.099 & 16.98 & ? & 10/03/2014 & 1.3 & 315 & 3\,$\times$\,1700.0 & 3\,$\times$\,1700.0 & 3.9 & \\ B11 & 187.959 & 58.35786 & 0.103 & 17.03 & P & 03/05/2014 & 1.4 & 446 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 4.2 & 2 \\ B12 & 190.78575 & 1.728797 & 0.092 & 17.09 & E & 05/05/2014 & 1.2 & 238 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 3.6 & \\ B13* & 191.979 & -3.627378 & 0.09 & 16.59 & S & 03/05/2013 & 0.7 & PA & 2\,$\times$\,1200.0 & 4\,$\times$\,900.0 & 2.3 & 2 \\ B14* & 192.3075 & 15.252789 & 0.083 & 16.90 & S & 01/05/2013 & 0.7 & PA & 4\,$\times$\,900.0 & 3\,$\times$\,900.0 & 2.3 & \\ B15* & 205.55083 & -0.293453 & 0.086 & 17.17 & E & 02/05/2013 & 0.7 & PA & 3\,$\times$\,900.0 & 4\,$\times$\,900.0 & 2.3 & \\ B16 & 207.66092 & 53.73111 & 0.108 & 16.95 & E & 09/03/2014 & 1.4 & 267 & 3\,$\times$\,1700.0 & 3\,$\times$\,1700.0 & 4.2 & \\ B17 & 211.27605 & 2.771761 & 0.077 & 17.17 & P & 04/05/2014 & 1.2 & 180 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 3.6 & 2 \\ B18 & 212.88733 & 45.28614 & 0.071 & 17.14 & E & 02/05/2014 & 1.6 & 109 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 6.0 & \\ B19* & 230.6967 & 59.35285 & 0.076 & 17.09 & P & 01/05/2013 & 0.8 & PA & 4\,$\times$\,1200.0 & 4\,$\times$\,900.0 & 2.6 & \\ B20!* & 231.55424 & 3.884864 & 0.086 & 16.79 & E & 02/05/2013 & 0.9 & PA & 3\,$\times$\,1200.0 & 3\,$\times$\,900.0 & 3.0 & \\ B21 & 234.29971 & 41.0717 & 0.098 & 16.68 & E & 28/07/2014 & 1.6 & 136 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 6.0 & \\ B22! & 245.43016 & 29.725689 & 0.098 & 16.50 & E & 29/07/2014 & 1.2 & 264 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 3.6 & \\ B23 & 327.73575 & -6.819708 & 0.059 & 16.68 & E & 26/07/2014 & 1.0 & 151 & 3\,$\times$\,1800.0 & 3\,$\times$\,1800.0 & 3.0 & \\ \hline \end{tabular} \end{center} \begin{flushleft} {\textbf{Column description:} \textbf{ID} - MLLINER identification (sources observed with NOT are marked with '*'; sources marked with '!' are possibly Sy2 galaxies and not LINERs as explained in Section~\ref{sec_emission_measure}); \textbf{RA}, \textbf{DEC} - J2000 right ascension and declination in degrees; \textbf{z} - redshift, from SDSS public catalogues; \textbf{m6500} - AB continuum magnitude at 6500\,\AA; \textbf{morph} - visual morphological classification where E, S, and P stand for Elliptical/S0, spiral, and peculiar (see the text); \textbf{Date} - date of observation; \textbf{seeing} - average FWHM of the seeing in arc-sec; \textbf{pos. ang.} - slit position angle in degrees (PA means that the paralactic angle was used, otherwise the angle is orientated along the major axis); \textbf{texp\_b} and \textbf{texp\_r} - total exposure time in blue and red parts in seconds; \textbf{Area$_{nuc}$} - area covered with our 'nuclear' extraction, in arc-sec$^2$ (just for comparison, the SDSS spectra cover an area of 7.08\,arc-sec$^2$); \textbf{IR data} - availability of Herschel (1) and IRAS (2) data.} \end{flushleft} \end{table*} \begin{figure*} \centering \includegraphics[width=0.8\textwidth,angle=0]{Observed_SDSSstamps_0.2arcsecperpix_part1_scaled.ps} \caption{SDSS gri colour images of our selected sample of the most luminous local LINERs. The top and bottom identifications correspond to our and SDSS ones, respectively. \label{fig_thumbnails_part1}} \end{figure*} \renewcommand{\thefigure}{\arabic{figure} (Cont.)} \addtocounter{figure}{-1} \begin{figure*} \centering \includegraphics[width=0.8\textwidth,angle=0]{Observed_SDSSstamps_0.2arcsecperpix_part2_scaled.ps} \caption{ \label{fig_thumbnails_part2}} \end{figure*} \renewcommand{\thefigure}{\arabic{figure}}
\label{sec_results_discussion} \subsection{General properties of the MLLINERs} \label{sec_discussion_general_properties} \indent In this section we describe the general properties of our MLLINERs: their masses, extinction, morphology, SFRs, and stellar populations. We compare them with the properties of other LLINERs (see Section~\ref{sec_sample_selection}), with the sample of the most-luminous LINERs at z\,$\sim$\,0.3 \citep{tommasin12}, and with the nearby and local LINER population analysed in previous studies \citep[e.g,][]{ho97,leslie16}. \subsubsection{Stellar and black hole mass} \label{sec_discussion_masses} \indent The nuclear stellar masses of our MLLINERs cover the range between 5.7\,$\times$\,10$^9$\,M$_{\odot}$ and 8.32\,$\times$\,10$^{10}$\,M$_{\odot}$. The median stellar mass is 1.52\,$\times$\,10$^{10}$\,M$_{\odot}$ and the average mass is 2.11\,$\times$\,10$^{10}$\,M$_{\odot}$. Fig.~\ref{fig_masses} (top plot) shows the distribution of our nuclear measurements. Using the SDSS MPA-JHU DR7 measurements of total stellar masses, we found that our MLLINERs cover the range 7.21\,$\times$\,10$^9$\,M$_{\odot}$\,-\,2.71\,$\times$\,10$^{11}$\,M$_{\odot}$, with median masses of 6.58\,$\times$\,10$^{10}$\,M$_{\odot}$. In Fig.~\ref{fig_masses} (bottom plot) we compared this distribution with those of LLINERs (see Fig.~\ref{fig_sample_selection}), and with the sample of the most luminous LINERs at z\,$\sim$\,0.3 from \cite{tommasin12}. Interestingly, MLLINERs at z\,$\sim$\,0.07 and z\,$\sim$\,0.3, although hosted by massive galaxies, do not cover the region of the most massive galaxies. When comparing our sample with the sample at z\,$\sim$\,0.3 the distributions are not completely consistent (Kolmogorov-Smirnov (KS) probability factor of 0.02). A significant part (35\%) of \cite{tommasin12} LINERs have lower stellar masses, however the peak of the two distributions at log\,M$_*$\,$\sim$\,10.9\,M$_{\odot}$ is the same for both samples.\\ \indent We compared the distributions of black hole masses (MBH) between MLLINERs and LLINERs. To derive MBH, we used its correlation with stellar velocity dispersion found by \cite{tremaine02} in the nearby universe, shown to be reliable for elliptical and bulge-dominated galaxies. We recovered stellar velocity dispersions from the MPA-JHU DR7 catalogue, and we obtained MBH only for galaxies classified as ellipticals (see Section~\ref{sec_observations}). The values are given in Table~\ref{tab_SFR_AGNluminosity}. MLLINERs cover the range between log\,(MBH/M$_{\odot}$)\,=\,7.03\,-\,8.57 with a median value of log\,(MBH/M$_{\odot}$)\,=\,7.45, while LLINERs show MBH in the range log\,(MBH/M$_{\odot}$)\,=\,6.24\,-\,8.54 and median value of log\,(MBH/M$_{\odot}$)\,=\,8.04. Interestingly, our MLLINERs do not contain the most massive BHs in their centres. \\ \begin{figure} \centering \begin{minipage}[c]{0.33\textwidth} \includegraphics[width=6.0cm,angle=0]{PAPER_IDL_histo_MstarNuclear_OurMeasurements_noSyB20noB22.ps} \end{minipage} \begin{minipage}[c]{0.33\textwidth} \includegraphics[width=6.0cm,angle=0]{PAPER_IDL_histo_MstarTotalSDSS_OurSample_noSyB20noB22_AllL2sample_Tommasin.ps} \end{minipage} \caption[ ]{\textit{Top:} Distribution of \textbf{nuclear} stellar masses of MLLINERs. \textit{Bottom:} Distributions of the SDSS/DR7 \textbf{total} stellar masses of MLLINERs (filled blue histogram), of the entire population of LLINERs (solid black lines), and of the most-luminous LINERs at z\,$\sim$\,0.3 (red dashed lines) from \cite{tommasin12}.} \label{fig_masses} \end{figure} \indent It could be surprising that MLLINERs, having on average higher SFRs than LLINERs, show in general lower stellar masses. Different works, both observational \citep{kauffmann03c,mateus06,leauthaud12,perez13} and numerical \citep{shankar06,behroozi12}, revealed a stellar mass of $\sim$\,6\,$\times$\,10$^{10}$\,M$_{\odot}$ as critical for the growth rate of stellar populations. In particular, in \cite{perez13} by studying a 3D spectroscopic sample of 105 local galaxies, the authors found that in galaxies more massive than 5\,$\times$\,10$^{10}$\,M$_{\odot}$ the inner regions ($<$\,0.5R$_{50}$) grew as much as 50\%\,-\,100\% faster than in the lower-mass galaxies. They found that the peak of relative growth rates of inner and outer galaxy regions correspond to the stellar mass of 6\,-\,7\,$\times$\,10$^{10}$\,M$_{\odot}$ (see their Fig. 5), while for lower and higher masses the growth rate decreases and therefore SFRs (LSF) as well. The median stellar mass of our MLLINERs (6.58\,$\times$\,10$^{10}$\,M$_{\odot}$) corresponds perfectly to this region, while for most LLINERs their stellar masses are already higher and correspond to lower values of the relative growth rate (lower LSF). This explains why MLLINERs having in average lower stellar masses in comparison to LLINERs, have higher LSF. \\ \indent In addition, we studied the stellar mass distributions of MLLINERs and LLINERs for the three morphological groups. Table~\ref{tab_masses_morphology} shows the median stellar masses for different morphological types. As can be seen, of three morphological types the highest difference was obtained for early-type galaxies. These galaxies represent a significant fraction of MLLINERs (40\%, see Fig.~\ref{fig_comparison_morphology}) and their median mass corresponds exactly to the highest relative growth rate of stellar populations (according to \cite{perez13}), which then could explain their high SFR values. This is not the case for early-type LLINERs that are characterised by higher stellar masses (lower growth rates) and lower SFRs. \begin{table} \small \begin{center} \caption{Median total stellar masses of MLLINERs and LLINERs in relation to morphology (given as logarithm and in M$_{\odot}$) \label{tab_masses_morphology}} \begin{tabular}{| c | c | c |c | c | c | c |} \hline &&\textbf{All}&\textbf{E}&\textbf{S}&\textbf{P}&\textbf{unclass}\\ \hline \hline &\textbf{logM$_{tot}$}&10.83&10.73&11.0&11.08&10.69\\ \hline \textbf{MLLINERs}&\textbf{st.dev.}&0.34&0.34&0.33&0.22&0.22\\ \hline &\textbf{num.}&40&16&8&10&6\\ \hline \hline &\textbf{logM$_{tot}$}&11.04&11.14&11.05&11.14&10.94\\ \hline \textbf{LLINERs}&\textbf{st.dev.}&0.33&0.36&0.29&0.35&0.31\\ \hline &\textbf{num.}&88&33&27&10&18\\ \hline \end{tabular} \end{center} \end{table} \subsubsection{Extinction} \indent We found that our MLLINERs can be hosted by galaxies with a wide range of extinctions. When using the Av measurements based on the H$\alpha$ and H$\beta$ emission lines, we find that most of them reside in galaxies with high extinctions. The median Av is 1.65\,mag, covering the range 0.49\,-\,3.46\,mag (see Fig.~\ref{fig_extinction}, top plot). For comparison, the bottom plot in Fig.~\ref{fig_extinction} shows the Av distributions of LLINERs and MLLINERs when taking into account the SDSS MPA-JHU DR7 3\,arcsec fibre measurements, where we measured Av in the same way as explained in Section~\ref{sec_emission_measure}, through H$\alpha$ and H$\beta$ lines. Both samples cover similar range of extinctions, having the majority of galaxies with higher values of Av\,$>$\,1.0. These values are higher than the extinctions of the nearby and low-luminosity LINERs in \cite{ho97}, with median value of Av\,=\,0.97. 54\% and 78\% of all nearby LINERs in \cite{ho97} have Av parameter $<$\,1.0 and $<$\,1.5, respectively . These comparisons between the two samples of LINERs are consistent with the general finding that the typical extinction increases with SFR \citep[e.g.,][]{kauffmann03a}. \begin{figure} \centering \begin{minipage}[c]{0.33\textwidth} \includegraphics[width=6.0cm,angle=0]{PAPER_IDL_histo_OurSample_AvEmissionLines_noSyB20noB22.ps} \end{minipage} \begin{minipage}[c]{0.33\textwidth} \includegraphics[width=6.0cm,angle=0]{PAPER_IDL_histo_AvEmissionLines_OurSample_noSyB20noB22_WholeL2sample_withHo.ps} \end{minipage} \caption[ ]{\textit{Top:} Distribution of Av in magnitudes of MLLINERs measured from emission lines. \textit{Bottom:} Distributions of SDSS fibre Av in magnitudes measured through emission lines of: MLLINERs (filled blue histogram), the entire population of LLINERs (solid black lines), and \cite{ho97} sample of nearby LINERs (red dashed lines). } \label{fig_extinction} \end{figure} \subsubsection{Morphology} \indent The MLLINERs studied in this work are hosted by galaxies with all morphologies, as shown in table~\ref{tab_obs} (see Section~\ref{sec_opt_spec_data} for classification details). Fig.~\ref{fig_comparison_morphology} shows comparisons between our MLLINERs (top plot) and LLINERs (bottom plot). While the differences per morphological type between MLLINERs and LLINERs are not significant ($\sim$\,10\% at most), by selecting MLLINERs we are selecting more E in comparison to S types.\\ \indent To compare our results with the sample by \cite{tommasin12} at z\,$\sim$\,0.3, we obtained the visual morphological classification in a completely consistent way as in our case, using the same classifiers and the same morphological types. We used HST/ACS images from the COSMOS\footnote{http://cosmos.astro.caltech.edu/} survey \citep{scoville07}, but we previously worsen their resolution to map the same physical size of $\sim$\,2kpc as in the case of SDSS images, and to have therefore comparable classifications. The fractions of E, S, P and unclassified galaxies can be seen in Fig.~\ref{fig_comparison_morphology} for FIR detected sample (top plot) and the entire optically-selected sample (bottom). When comparing z\,$\sim$\,0.3 and our samples, it seems that the fraction of galaxies classified as peculiar is similar at both redshifts and in both plots. On the other hand, we find higher fraction ($\sim$\,20\%) of early-type galaxies in our samples and of spiral galaxies in \cite{tommasin12}. To confirm if the observed differences are significant, we need better statistics. Incompleteness of the sample at z\,$\sim$\,0.3, plus the selection effects could be responsible for the observed differences. The most luminous \cite{tommasin12} galaxies were selected in the FIR using Herschel data, while our MLLINERs selection was carried out in optical. This could be the reason for the differences observed in the top plot of Fig.~\ref{fig_comparison_morphology}. If we check the morphological classification of our MLLINERs with the available FIR data (Table~\ref{tab_obs_fir}), we also observe that most sources are later types (68\%), classified either as S or peculiar. The sample is again too small for providing any reliable conclusions. On the other side, spectroscopic classification methods applied on the entire \cite{tommasin12} sample also differs from ours, and were based on NII-BPT and/or SII-BPT diagrams, while we used NII-BPT and OI-BPT diagrams (see sec.~\ref{sec_sample_selection}). Moreover, the apertures used in our and in \cite{tommasin12} samples cover different physical sizes of the observed galaxies. \\ \indent Different criteria were used in \cite{ho08} and this work to classify galaxies morphologically. While we are dealing with low-resolution data (and therefore only a rough classification in three morphology groups, E, S and P, is made) Ho's sample of nearby LINERs provides very detailed information on morphological structures. Therefore, since we are not dealing with samples classified in a consistent way, we are not able to provide any direct comparison with Ho's sample. In general, we would like to stress that our MLLINERs show higher fractions of later-types in comparison to nearby LINERs. Moreover, a significant fraction ($\sim$\,25\%) of MLLINERs are hosted by peculiar systems, showing unusual structures and clear signs of interactions, at both low- and higher-redshifts, which is again in contrast with the morphology of nearby LINERs. \begin{figure} \centering \begin{minipage}[c]{0.33\textwidth} \includegraphics[width=6.0cm,angle=0]{PAPER_IDL_morphology_OurSample_noSYB20B22_TommasinFIRsample_percentage_gimp.ps} \end{minipage} \begin{minipage}[c]{0.33\textwidth} \includegraphics[width=6.0cm,angle=0]{PAPER_IDL_morphology_allL2Sample_entireTommasinsample_percentage_gimp.ps} \end{minipage} \caption[ ]{\textit{Top:} Fraction of MLLINERs per morphological type (red filled circles)in our sample and in the \cite{tommasin12} FIR Herschel sample (blue filled circles). \textit{Bottom:} Entire LLINER sample and the entire \cite{tommasin12} optically selected sample (blue filled circles). E, S, P, and Unclass stand for Ell/S0, spiral, peculiar, and unclassified galaxies, respectively (see sec.~\ref{sec_opt_spec_data})} \label{fig_comparison_morphology} \end{figure} \subsubsection{SFRs} \indent In this work we use three different measurements of SFRs (see Section~\ref{sec_sfr}), two based on optical data (spectral fitting and Dn4000 index) and one on FIR (Herschel and IRAS). The average nuclear SFRs measured with STARLIGHT and Dn4000 index is $\sim$\,3\,[M$_{\odot}$/yr], which is significantly smaller than the SFR inferred from FIR observations with an average of $\sim$\,13\,[M$_{\odot}$/yr]. Most of the difference must be due to the fact that the nuclear region, in all sources, is considerably smaller than the size of the galaxy. If we scale the optical measurements of SFRs to the entire galaxy, assuming that the sSFR is constant (see Sec.~\ref{sec_sfr}), the difference between the optical and FIR methods becomes smaller: the average SFRs in this case are $\sim$\,9\,[M$_{\odot}$/yr] and $\sim$\,11\,[M$_{\odot}$/yr] when using STARLIGHT best-fit models and Dn4000 index, respectively. \\ \indent Fig.~\ref{fig_SFR_comparison} shows the comparison between SFRs obtained through different methods. With two different and independent methods based on optical data (spectral fitting and strength of 4000\,\AA\,Balmer break), we obtained consistent measurements of SFR, as can be seen on the top plot. As explained in Section~\ref{sec_sfr}, we used simulations from \cite{brinchmann04} to extract the mode sSFR for our nuclear measurements of Dn4000. This could be a source of several uncertainties. First, we are using just the mode values while for each Dn4000 the range of possibilities is much wider. In addition, Dn4000 measurements are based on nuclear spectra in this work while the authors used the information from SDSS aperture which is larger (see table~\ref{tab_obs} and Section~\ref{sec_observations}). Finally, in this work we are dealing with MLLINERs while the simulations were done for star-forming galaxies. Despite all this, we find a good agreement between STARLIGHT and Dn4000 SFR measurements, with $\sim$\,90\% of the sample being inside a difference of 1\,$\sigma$. \\ \indent When comparing the FIR estimations with the optical ones, but scaled to match the entire galaxy, the dispersion is larger, as shown in Fig.~\ref{fig_SFR_comparison} (bottom plot). We found $\sim$\,50\% of the sample with differences higher than 1\,$\sigma$, however we don't see any systematic trend. Several possibilities can explain the differences. First, as mentioned above we are dealing with different apertures, not only when comparing optical and FIR estimations, but for Herschel and IRAS. Secondly, the scaling assumed here, that the sSRF for the slit and the entire galaxy is the same, can lead to large uncertainties. There are other possibilities related to the geometry of the obscuring dust that affect the optically-based method much more than the FIR-based methods.\\ \indent Discrepancies based on optical and FIR SFR measurements were reported in previous works, usually finding smaller optical values in comparison to FIR \citep[e.g.][]{rigopoulou00,cardiel03,wuyts11,tommasin12}, but the scatter in most of these works is larger than in our case. In sample of the most luminous LINERs at z\,$\sim$\,0.3 by \cite{tommasin12}, the H$\alpha$ and UV measurements of SFRs are $\sim$\,30 times smaller than the FIR measurements. In contrast, the typical FIR SFRs in their sample are $\sim$\,10\,[M$_{\odot}$/yr], similar to our FIR estimations. \begin{figure} \centering \begin{minipage}[c]{0.43\textwidth} \includegraphics[width=6.0cm,angle=0]{IDLplots_SFRscomparisons_SFRstarlight_vs_SFRDn4000_FINAL_withERRORS_2_noB20noB22.ps} \end{minipage} \begin{minipage}[c]{0.43\textwidth} \includegraphics[width=6.0cm,angle=0]{IDLplots_SFRscomparisons_SFRDn4000scaled_vs_SFRfir_FINAL_withERRORS_2_noB20noB22.ps} \end{minipage} \caption[ ]{\textit{(From top to bottom)} Comparison between SFRs measured with different methods: STARLIGHT and Dn4000 index for nuclear spectra, and Dn4000 (scaled) and FIR data. FIR data contain information from both Herschel\,-\,PACS (filled circles) and IRAS (open triangles). } \label{fig_SFR_comparison} \end{figure} \subsubsection{Stellar populations and star formation histories} \indent As shown in Section~\ref{sec_starlight_fits} the nuclear regions of our MLLINERs are mainly characterised by intermediate (10$^8$\,$<$\,age\,[yr]\,$\le$\,10$^9$) and old (age\,[yr]\,$>$\,10$^9$) stellar populations. In $\sim$\,30\% of the sources the contribution of both intermediate and old stars is similar. In $\sim$\,20\% and 45\% of MLLINERs intermediate and old stellar populations are dominant, respectively. A young (age\,[yr]\,$\le$\,10$^8$) stars population is found in the nuclear regions of our sources in 43\% of MLLINERs, but for most of these galaxies the young stellar populations represent only $<$\,10\% of all stars. The median age of MLLINERs is logt\,=\,8.97\,[yr], covering the range logt\,=\,8.17\,-\,9.82\,[yr]. Our results are consistent with previous findings for low-luminous AGN (LINERs included) whose nuclear regions contain intermediate and old stellar populations \citep{cid04,gonzalez04}. Most of SF measured in FIR is possibly related to circumnuclear regions of MLLINERs, due to high stellar masses and/or young stars, since with our nuclear spectra in average we only cover $\sim$\,30\% of the total stellar mass. \\ \indent As mentioned in Section~\ref{sec_dn4000_hdelta}, the Dn4000 and H$\delta$ indices can be used as indicators of the SFH. The location of galaxies in the Dn4000 vs. H$\delta$ diagram, shown to be a powerful diagnostic of whether they have been forming stars continuously or in bursts over the past 1\,-\,2\,Gyr. Galaxies with continuous SFHs occupy a narrow strip in this plane (see Fig.~\ref{fig_Dn4000_Hdelta}). Following \cite{kauffmann03} models (Fig.~\ref{fig_Dn4000_Hdelta}), twelve MLLINERs (F12, F17, F21, F23, F24, B02, B15, B16, B17, B18, B19, and B23) might have experienced a burst of SF over more than 0.1\,Gyr ago (green circles). One source (B13) possibly experienced a burst of SF over less than 0.1\,Gyr ago (yellow circle). The other 17 sources (F02, F03, F06, F09, F14, F15, F20, F22, B01, B04, B05, B06, B08, B10, B11, and B14) could suffer both, burst and continuous SF (red circles). F04 has F$_{burst}$\,=\,0 (orange circle), and one can say with high confidence that this galaxy did not form a significant fraction of its stellar population in a burst over a past 2\,Gyr. Finally, five MLLINERs (F01, F16, F19, B07, and B12) lie outside the range covered by models (violet circles). \subsection{AGN and SF luminosities of MLLINERs} \label{sec_discussion_tommasinRelation} \indent The connection between LSF and LAGN was studied in many previous works, at different redshifts and for different samples of AGN, leading to somewhat inconsistent results \citep[e.g.,][and references therein]{netzer09,lutz10,rovilos12,page12,santini12,barger15,azadi15,netzer16}. Such relationships have been studied for AGN dominated sources (LAGN\,$>$\,LSF), SF dominated sources (LSF\,$>$\,LAGN) and the entire population. Some of the suggested correlations are clearly related to the sample selection (e.g., FIR or X-rays) and averaging (e.g., stacking) methods. In this section we study the relationship at low redshift for our samples of MLLINERs and LLINERs. Figure~\ref{fig_tommasin} shows LSF vs. LAGN for our two samples, where MLLINERs are represented with coloured filled circles and LLINERs with black dots. For MLLINERs, LAGN and LSF were measured as explained in previous sections. In the case of LLINERs, we used the SDSS/DR7 data and applied the H$\beta$ and OIII+OI methods to obtain LAGN, and the scaled Dn4000 method to obtain LSF. We note that in this case, some of the measured Dn4000 indices are very large (1.7 or larger) and hence cannot be used to obtained reliable SFRs \citep{kauffmann03}. We estimate this threshold to be equivalent to $\sim$\,logLSF\,=\,42.9 erg/sec (about 0.2\,M$_{\odot}$/yr).\\ \indent Figure~\ref{fig_tommasin} shows that MLLINERs tend to lie on the one-to-one LSF-LAGN relation (indicated on the diagram with a dashed line). About 90\% of all MLLINERs have values of LSF and LAGN in the range 10$^{44}$\,-\,10$^{45}$\,erg/sec. For comparison, we plotted also the line indicating the location of AGN-dominated galaxies from \cite{netzer09} (dotted line) which, by definition, are located below the 1:1 line. Our MLLINERs are located clearly above this line, and remain closer to the 1:1 relationship. On the other side, LLINERs are located below the one-to-one LSF-LAGN line, showing a wide range of LSF for the same LAGN. We suggest that this is again related to the stellar mass differences between MLLINER and LLINER samples discussed in sec.~\ref{sec_discussion_masses}. Although having the same LAGN, LLINERs with stellar masses higher than the critical one (of 6\,-\,7\,$\times$\,10$^{10}$\,M$_{\odot}$) seem to have already lower relative growth rates of stellar populations, and therefore lower LSF. As shown in \cite{perez13}, the differences in the growth rate can be even 50\%\,-\,100\%, which could explain significant differences in LSF between LLINERs and MLLINERs for the same LAGN.\\ \indent We compared our results with those for the most-luminous LINERs at z\,$\sim$\,0.3 using again the sample of \cite{tommasin12}. We used their measurements of LAGN and LSF, where LAGN were derived from the H$\beta$ and O[III]$\lambda$5007 methods and LSF from Herschel observations. In general, the location of MLLINERs at z\,$\sim$\,0.04\,-\,0.11 and at z\,$\sim$\,0.3 are very similar. \cite{tommasin12} compared their results with nearby LINERs from \cite{ho97}, finding that the later are characterised by considerably lower LAGN and LSF. In Fig.~\ref{fig_tommasin} we marked the region that corresponds to the location of nearby LINERs (blue dashed box). As can be seen, although both AGN and SF luminosities show lower values, in this case the dispersion from 1:1 relation is much larger. While some sources are distributed around 1:1 relation the others lie more around the AGN-dominated line. Note also that for low LAGN ($\sim$\,10$^{41}$\,erg/s) the difference between 1:1 and AGN-dominated relations becomes less significant \citep{netzer09}. \\ \indent In order to explain the differences between nearby and z\,$\sim$\,0.3 LINERs, \cite{tommasin12} pointed out several possibilities. First, the aperture difference, which is much smaller in the case of the Ho's sample, where only the very central regions of the galaxies are included. Second, the FIR selection of the z\,$\sim$\,0.3 sample in comparison to the Ho's LINERs, enforces higher values of LSF. Third, they argue that LINERs with such high LSF could be present in the local universe, but have not been studied yet systematically. Finally, \cite{tommasin12} suggested that there might be a real evolution in AGN and SF luminosities between z\,$\sim$\,0 and z\,$\sim$\,0.3. With our work we can provide more information about some of the questions made by \cite{tommasin12}. We can confirm the existence of LINERs in the local universe with the same SF and AGN properties as at z\,$\sim$\,0.3, discarding therefore the pure evolutionary scenario.\\ \begin{figure} \centering \includegraphics[width=0.4\textwidth,angle=0]{IDLplots_TommasinRelation_logLagn_vs_logLSF_FINAL_zoomed_withERRORS_noSyB20noB22_HoInBox.ps} \caption{The relationship between the AGN and SF luminosities of the most luminous local LINERs. LSF was measured in three different ways: with Herschel/PACS FIR data (big green filled circles), IRAS data (big dark blue filled circles), and through Dn4000 index (big red filled circles). For comparison, we plot the entire sample of LLINERs (small black dots), and Tommasin et al. (2012) sample of the most luminous LINERs at z\,$\sim$\,0.3 (black crosses). The blue dashed box shows the area where the nearby LINERs from Ho et al. (1997) are located. The dashed line shows the one-to-one LAGN-LSF relation, while the dotted line shows the empirical relationship for AGN-dominated sources from Netzer et al. (2009). The horizontal dashed-dot-dashed line shows the limit below which we do not trust LSF (at about 8\,$\times$\,10$^{42}$\,=\,0.2\,M$_{\odot}$/yr) \label{fig_tommasin}} \end{figure} \begin{figure} \centering \includegraphics[width=0.4\textwidth,angle=0]{IDLplots_TommasinRelation_StanleyAGNz1_OurSample3LumBinsALL_TommasinSample2bins_RealMeansWithAllSamples_FINAL.ps} \caption{The relationship between AGN and SF luminosities for all luminous LINERs, divided in three different LAGN bins (black, green, and blue crosses). The average values of LSF and LAGN per bin are represented with black filled circles. For comparison we show the average values of Stanley et al. (2015) for X-ray detected AGN in their first redshift bin of z\,$\sim$\,0.4 (red diamonds) and the average of the entire sample of Tommasin et al. (2012) at z\,$\sim$\,0.3 (blue triangles). \label{fig_stanley}} \end{figure} \indent Recently, \cite{stanley15} studied the relationship between LSF and LAGN for a sample of $\sim$\,2000 X-ray detected (10$^{42}$\,$<$\,L$_{2-8keV}$\,$<$\,10$^{45.5}$\,erg/sec) AGN at redshifts z\,=\,0.2\,-\,2.5. They divided all galaxies in four redshift ranges, and for each redshift range they measured the average LSF and LAGN in bins of 40 galaxies. LSF was measured using FIR data, and was based mostly on Herschel upper limits (which is why they could only discuss mean LSF). LAGN is based on X-ray 2\,-8\,keV measurements. They found that the relationship between the average LSF and LAGN is mainly flat, independently of redshift and AGN luminosity. To test the flatness of the observed relationship, the authors tested their results with two empirical models \citep{aird13, hickox14} that predict $<$\,LSF\,$>$ as a function of LAGN. They suggested that the flat relationship is due to short-time scale variations in LAGN caused by changes in mass accretion rate onto the BH. These variations are shorter than those related to SF, and therefore for a given value of mean LSF, AGN luminosity can take different values and flatten the correlation. \\ \indent Here we are able to test, for the first time, \cite{stanley15} results for LINERs. We used the entire sample of LLINERs (MLLINERs included), dividing it in three LAGN bins (with 43 galaxies in the first bin and 44 in the other two bins) and measured mean LAGN and LSF in each bin. Figure~\ref{fig_stanley} shows all sources with crosses, while the mean LAGN and LSF values in the three LAGN bins are marked with filled black circles. For comparison, we plotted \cite{stanley15} averaged values for their first redshift bin at z\,$\sim$\,0.4 (red diamonds). We also show the results for the LINERs in \cite{tommasin12}. Only 34 out of the 97 objects in the Tommasin sample have measured (Herschel) SFRs. Since we are comparing averaged properties, we assume that all other LINERs in that sample have LSF\,=\,0. This would mean that the numbers we use are somewhat smaller than the actual mean LSF. \\ \indent Our results considering LLINERs are in general agreement with the \cite{stanley15} results. However, we do not have to rely on mean properties and can look at the entire LSF distribution in each bin of LAGN. The measured range in LSF is large, about 1.5 dex, similar to the overall range in LAGN. Obviously, using mean values will tend to emphasize the larger number of low SFR sources in each bin. However, the sources with the highest LSF in each LAGN bin certainly have different properties than the ones with the lowest LSF, as discussed in the following section. \subsection{MLLINERs and the main sequence of SF galaxies} \label{sec_discussion_ms} \indent SF galaxies show a tight and well-defined relationship called the 'main sequence' (MS) between SFR and stellar mass. This relationship depends on redshift and has been studied at different cosmic time \citep[e.g.][and references therein]{brinchmann04,noeske07,elbaz07,daddi07,gonzalez10,whitaker12,guo13,leslie16}. Figure~\ref{fig_ms} shows all the objects studied in this work on the SFR\,-\,M$_*$ diagram. For the SFRs we used exactly the same data as in Fig.~\ref{fig_tommasin}. For the stellar mass we used the mass of the entire galaxy, recovered from the MPA-JHU DR7 catalogue. For the MS, we used the fit obtained by \cite{whitaker12}, whose SFRs are also based on Kroupa IMF. We plotted the MS (solid line) for z\,=\,0.07, which is the average value in our sample. For the width of the MS we used $\pm$\,0.3\,dex (dashed lines), found in many previous works to be the typical 1\,$\sigma$ boundaries \citep[e.g,][]{elbaz07,rodighiero10,whitaker12,whitaker14,shimizu15}. More than 90\% of our MLLINERs lie along the main sequence of SF galaxies (within the dashed lines). \begin{figure} \centering \includegraphics[width=0.4\textwidth,angle=0]{IDLplots_SFmainsequence_withWidth_FINAL_withERRORS_2_withLeslie_noSyB20noB22_DiffSizeSymbols_LLINERsalso.ps} \caption{The relationship between SFR and total stellar mass. SFRs were measured in three different ways: with Herschel/PACS FIR data (big green filled circles), IRAS data (big dark blue filled circles), and through Dn4000 index (big red filled circles). The solid black line shows the Whitaker et al. (2012) fit for the main sequence, and the dashed lines its typical width (see the text). The entire sample of luminous LINERs (small black dots), and Tommasin et al. (2012) sample of the most luminous LINERs at z\,$\sim$\,0.3 (black crosses) are shown for comparison. The dotted area is reproduced from Leslie et al (2016) and represents the typical location of 60\% of all LINERs at low redshifts. Depending on their AGN luminosity, MLLINERs and LLINERs are represented with symbols of different sizes (see sec. \ref{sec_discussion_ms_LAGN}). \label{fig_ms}} \end{figure} \indent Once again our MLLINERs at z\,=\,0.04\,-\,0.11 show the same properties as the most luminous LINERs at z\,$\sim$\,0.3 (black crosses in Fig.~\ref{fig_ms}) in \cite{tommasin12}. At both redshifts, the most luminous LINERs represent $\sim$\,1/3 of all LLINER. Most remaining 2/3 of LLINERs lie below the MS (black dots), having lower SFRs for masses typical of MLLINERs or even higher. Considering morphological types, we found that the different types are located on the MS. This sample seems to be different from the general galaxy population where later-types are mainly located on the MS, while earlier-types lie below it \citep[e.g.][and references therein]{gonzalezdelgado16}. \\ \indent Recently, \cite{leslie16} studied the SFR\,-\,stellar mass plane for different types of low-redshift galaxies from the SDSS survey. They classified all galaxies into star-forming, composite, Sy2, LINERs\footnote{Note that this work does not take into account the separation of LINERs into systems excited by AGN and by pAGB stars.}, and ambiguous, using the emission line ratios from MPA-JHU DR7 catalogues. 6.5\% of of the sources studied in this work are LINERs. We assumed this sample (of 13,176 galaxies) to be representative of LINERs at low redshifts and plot in Fig.~\ref{fig_ms} a dotted box representing $>$\,60\% of the sources in \cite{leslie16}. The average stellar masses and SFRs they found are $<$\,log(M$_*$)\,$>$\,=\,10.74 and $<$\,log(SFR)\,$>$\,=\,-0.79, respectively. These values are smaller than for our MLLINERs, $<$\,log(M$_*$)\,$>$\,=\,10.82 and $<$\,log(SFR)\,$>$\,=\,0.86, respectively. This is not surprising given that our MLLINERs were selected according to both LAGN and LSF. \subsubsection{\textbf{Relation between the fraction of SF galaxies and AGN luminosity}} \label{sec_discussion_ms_LAGN} The more important issue of the location of LINERs in the SFR vs. M$_*$ plane as a function of LAGN, as found here, was not considered by \cite{leslie16}. To illustrate this we consider the properties of all SDSS/DR7 LINERs in the redshift range 0.04\,-\,0.11. We measured LAGN as described above and used the scaled Dn4000 method to estimate LSF. We then estimated their fraction on the MS using different bins of LAGN, where the MS is defined exactly as in Fig.~\ref{fig_ms}. The fraction of z\,=\,0.04\,-\,1.11 LINERs located on the MS is 2\%, 3\%, 11\%, and 37\% in the bins of logLAGN\,=\,43\,-\,43.5, 43.5\,-\,44, 44\,-\,44.5, 44.5\,-\,45, respectively. Thus we can safely conclude that the fraction of SF galaxies among low redshift LINERs is LAGN-dependent. While studies like those of Leslie et al. (2016) are not available at higher redshifts, it seems that for the most luminous LINERs, this difference from the rest of the population extends at least to z\,=\,0.3.\\ \indent In this work we analyse the properties of the 42 most-luminous LINERs (in terms of AGN and star-formation luminosities) at z\,=\,0.04\,-\,0.11 from the entire SDSS DR4 survey. We obtained long-slit spectroscopy of the nuclear regions for all sources, and FIR data (Herschel and IRAS) for 30\% of the sample. We carried out spectral fitting using the STARLIGHT code and templates from \cite{bruzual03}, testing 25 ages and solar metallicity. From the best-fit models we obtained the emission spectra, stellar masses, SFRs, stellar populations, and ages. We used the spectra to measure the emission lines, extinction, and extinction corrected luminosities. We also measured the Dn4000 and H$\delta$ indices. The AGN luminosities were measured through extinction-corrected emission lines, and SFRs using different indicators (both optical and FIR).\\ \indent Previous works characterised the population of local LINERs as: hosted by old and massive early-type galaxies, with low extinctions, massive black holes, old stellar populations and little star-formation \citep{ho97,ho08,heckman14}. In contrast, our most-luminous LINERs are hosted by both early- and late-types. Moreover, $\sim$\,25\% of sources are peculiar systems, with clear signs of sub-structures and interactions or mergers. We found higher values of extinction than typical for most low-redshift LINERs. The nuclear regions mainly consist of intermediate (10$^8$\,$<$\,age\,[yr]\,$\le$\,10$^9$) and old (age\,[yr]\,$>$\,10$^9$) stellar populations, while young stars are present only in 43\% of sources, similar to what has been found for nearby LINERs \citep{cid04}. The median SFRs are $\sim$\,10\,[M$_{\odot}$/yr], much higher than those for most local LINERs. However, it is interesting that they do not have the highest stellar masses, and in general show lower masses than other luminous LINERs. We found that the median stellar mass of our most-luminous LINERs corresponds to the mass of 6\,-\,7\,$\times$\,10$^{10}$\,M$_{\odot}$ measured in different works to be critical for the peak of relative growth rates of stellar populations (highest SFRs and LSF). Other LINERs although showing the same AGN luminosities, show lower SF luminosities. \\ \indent LINERs with these kind of properties were previously studied only at z\,$\sim$\,0.3 \citep{tommasin12}. With our work we confirmed the existence of such LINERs also at low-redshifts (z\,$\sim$\,0.07). They show the same properties in terms of stellar mass, SFRs, and AGN luminosity at both redshifts. Our most luminous LINERs tend to lie along the LAGN\,=\,LSF line hinting for co-evolution of the two properties. In addition, most of them are found on the MS of SF galaxies, with stellar masses $\gtrsim$\,10$^{10}$\,M$_{\odot}$. Finally, using the entire DR7 sample, we present evidence that the fraction of LINERs on the MS depends on their AGN luminosity.
16
7
1607.03915
This work presents the properties of 42 objects in the group of the most luminous, highest star formation rate (SFR) low-ionization nuclear emission-line regions (LINERs) at z = 0.04-0.11. We obtained long-slit spectroscopy of the nuclear regions for all sources, and FIR data (Herschel and IRAS) for 13 of them. We measured emission-line intensities, extinction, stellar populations, stellar masses, ages, active galactic nuclei (AGN) luminosities, and SFRs. We find considerable differences from other low-redshift LINERs, in terms of extinction, and general similarity to star-forming galaxies. We confirm the existence of such luminous LINERs in the local universe, after being previously detected at z ∼ 0.3 by Tommasin et al. The median stellar mass of these LINERs corresponds to 6-7 × 10<SUP>10</SUP> M<SUB>⊙</SUB> which was found in previous work to correspond to the peak of relative growth rate of stellar populations and therefore for the highest SFRs. Other LINERs although showing similar AGN luminosities have lower SFR. We find that most of these sources have LAGN ∼ LSF suggesting co-evolution of black hole and stellar mass. In general, the fraction of local LINERs on the main sequence of star-forming galaxies is related to their AGN luminosity.
false
[ "lower SFR", "similar AGN luminosities", "stellar populations", "stellar mass", "stellar masses", "local LINERs", "such luminous LINERs", "Other LINERs", "LINERs", "relative growth rate", "et al", "SFR", "SFRs", "active galactic nuclei", "black hole", "other low-redshift LINERs", "IRAS", "low-ionization nuclear emission-line regions", "previous work", "AGN" ]
15.07792
7.597906
119
2999276
[ "Walton, D. J.", "Middleton, M. J.", "Pinto, C.", "Fabian, A. C.", "Bachetti, M.", "Barret, D.", "Brightman, M.", "Fuerst, F.", "Harrison, F. A.", "Miller, J. M.", "Stern, D." ]
2016ApJ...826L..26W
[ "An Iron K Component to the Ultrafast Outflow in NGC 1313 X-1" ]
77
[ "Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA 91109, USA ; Space Radiation Laboratory, California Institute of Technology, Pasadena, CA 91125, USA", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "INAF/Osservatorio Astronomico di Cagliari, via della Scienza 5, I-09047 Selargius (CA), Italy", "Universite de Toulouse, UPS-OMP, IRAP, Toulouse, France ; CNRS; IRAP; 9 Av. colonel Roche, BP 44346, F-31028 Toulouse cedex 4, France", "Space Radiation Laboratory, California Institute of Technology, Pasadena, CA 91125, USA", "Space Radiation Laboratory, California Institute of Technology, Pasadena, CA 91125, USA", "Space Radiation Laboratory, California Institute of Technology, Pasadena, CA 91125, USA", "Department of Astronomy, University of Michigan, 1085 S. University Avenue, Ann Arbor, MI 49109-1107, USA", "Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA 91109, USA" ]
[ "2016ApJ...829...28B", "2017AN....338..234P", "2017ARA&A..55..303K", "2017ApJ...834...77F", "2017ApJ...839...46S", "2017ApJ...839..105W", "2017IJMPD..2630021M", "2017MNRAS.468.2865P", "2017MNRAS.469L..99F", "2017Natur.543...83P", "2017NewAR..79...26M", "2017RAA....17...63C", "2018ApJ...853..115W", "2018ApJ...856..128W", "2018ApJ...857L...3W", "2018ApJ...864L..27F", "2018MNRAS.473.4360W", "2018MNRAS.473.5680K", "2018MNRAS.475..154M", "2018MNRAS.475L.101D", "2018MNRAS.477.3711J", "2018MNRAS.478.5638M", "2018MNRAS.479.3978K", "2018MNRAS.480.4443T", "2018MNRAS.481..639J", "2018NatAs...2..312B", "2018SPIE10699E..29M", "2019ApJ...871..115Z", "2019ApJ...873...19T", "2019ApJ...877...57Q", "2019ApJ...883...44W", "2019ApJ...885L..38F", "2019MNRAS.483.5554E", "2019MNRAS.484.2544W", "2019MNRAS.486....2M", "2019MNRAS.489..282M", "2019MNRAS.489..366K", "2020ApJ...889...71B", "2020ApJ...895..127B", "2020Galax...8...70C", "2020MNRAS.491.3730K", "2020MNRAS.491.5172K", "2020MNRAS.491.5702P", "2020MNRAS.492.4646P", "2020MNRAS.493.2518F", "2020MNRAS.494.6012W", "2020NatAs...4..147B", "2021A&A...649A.104G", "2021A&A...651A..54M", "2021A&A...652A.118R", "2021ApJ...906...36Q", "2021ApJ...908..156Y", "2021ApJS..255....7R", "2021MNRAS.501.1002W", "2021MNRAS.501.1644S", "2021MNRAS.506.1045M", "2021MNRAS.508.3569K", "2022A&A...666A.100G", "2022ApJ...925..151T", "2022ApJ...929..138B", "2022ApJ...934...35M", "2022ApJ...935...38Z", "2022ApJ...938...76G", "2022MNRAS.509.1587W", "2022MNRAS.509.2493K", "2023A&A...669A.130A", "2023AN....34420134P", "2023ApJ...947...52Z", "2023ApJ...950..160L", "2023FrASS..1089432B", "2023MNRAS.519.2224M", "2023MNRAS.524.4302K", "2023NewAR..9601672K", "2023PASP..135e4101P", "2023arXiv230407977M", "2024A&A...682A..94B", "2024MNRAS.531..550K" ]
[ "astronomy" ]
9
[ "black hole physics", "X-rays: binaries", "X-rays: individual: NGC 1313 X-1", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Astrophysics of Galaxies" ]
[ "1996ASPC..101...17A", "1996ApJ...465..487V", "2000ApJ...542..914W", "2001A&A...365L...1J", "2001A&A...365L..18S", "2001A&A...365L..27T", "2001ApJS..133..221K", "2003MNRAS.345..705P", "2004ApJ...601..450M", "2004ApJS..154..519S", "2005A&A...431..111B", "2005A&A...440..775K", "2006ApJ...650L..75F", "2006MNRAS.368..397S", "2007A&A...464.1155C", "2007MNRAS.377.1187P", "2009MNRAS.393L..41K", "2009MNRAS.397.1836G", "2009MNRAS.400..677Z", "2010A&A...521A..57T", "2011ApJ...731L..32M", "2011MNRAS.413.1623D", "2011MNRAS.416.1844W", "2012ApJ...746L..20K", "2012MNRAS.420.1107P", "2012MNRAS.426..473W", "2013ARA&A..51..511K", "2013ApJ...770..103H", "2013ApJ...773L...9W", "2013ApJ...776L..36M", "2013ApJ...778..163B", "2013PASJ...65...88T", "2014ApJ...784L...2K", "2014ApJ...785L...7M", "2014ApJ...793...21W", "2014MNRAS.438L..51M", "2014Natur.514..202B", "2014SSRv..183..223C", "2015ApJ...799..121R", "2015ApJ...799..122W", "2015ApJ...806...65W", "2015ApJ...808...64M", "2015ApJS..220....8M", "2015MNRAS.447.3243M", "2015MNRAS.454.3134M", "2015Sci...347..860N", "2016ApJ...822L..18M", "2016ApJ...826...87W", "2016Natur.533...64P", "2016PhRvL.116f1102A" ]
[ "10.3847/2041-8205/826/2/L26", "10.48550/arXiv.1607.03124" ]
1607
1607.03124_arXiv.txt
Ultraluminous X-ray sources (ULXs) are variable, off-nuclear point sources in nearby galaxies with X-ray luminosities $L_{\rm{X}}\geq10^{39}$\,\ergps\ (\citealt{Swartz04, WaltonULXcat}). The brighter members of this population have luminosities that significantly (factors of 10 or more) exceed the Eddington limit for the $\sim$10\,\msun\ stellar-remnant black holes observed in accreting Galactic black hole binaries (\citealt{Casares14rev}). Multi-wavelength observations have largely ruled-out strong anisotropic emission as a means of skewing luminosity estimates (\citealt{Moon11}, although moderate collimation is still permitted). ULXs must therefore either host large black holes, potentially either the long-postulated yet elusive `intermediate mass' black holes ($M_{\rm{BH}}\sim10^{2-5}$\,\msun; \citealt{Miller04}) or the massive stellar remnants ($M_{\rm{BH}}\sim30-100$\,\msun; \citealt{Zampieri09}) recently confirmed by LIGO (\citealt{Abbott16gw}), or represent an exotic, highly super-Eddington accretion phase (\citealt{Poutanen07}). In either case, they hold clues to the processes governing the formation and evolution of supermassive black holes in the early Universe (\citealt{Kormendy13rev}). Since launch, the \nustar\ mission (\citealt{NUSTAR}) has undertaken a substantial program observing a sample of extreme ULXs, revealing the high-energy ($E>10$\,keV) behavior of these enigmatic sources for the first time. As one of the few sources within $\sim$5\,Mpc to persistently radiate at $L_{\rm{X}}\sim10^{40}$\,\ergps\ (\citealt{Miller13ulx}), \ngc\ ($D\sim4$\,Mpc) was an important part of this program, observed in coordination with \xmm\ (\citealt{XMM}) to provide broadband ($\sim$0.3--30\,\kev) spectral coverage. These observations revealed broadband spectra inconsistent with standard modes of sub-Eddington accretion (\citealt{Bachetti13, Miller14}; similar to other ULXs observed by \nustar\ to date, \citealt{Walton14hoIX, Walton15, Walton15hoII, Rana15, Mukherjee15}), supporting the idea that these sources are exhibiting a super-Eddington phase of accretion. Indeed, we now know at least one of these sources is a highly super-Eddington neutron star (\citealt{Bachetti14nat}). A prediction of all super-Eddington accretion models is that powerful winds should be launched (\citealt{Poutanen07, King09, Dotan11, Takeuchi13}). Robust detection of any such winds from ULXs has, however, proven challenging (\citealt{Walton12ulxFeK, Walton13hoIXfeK}). For \ngc, \cite{Middleton15soft} report low-energy ($\sim$1\,keV) blended atomic features that are consistent with being absorption from an ionised outflow, but the low-resolution CCD spectra considered prevented a conclusive identification as such. However, in a key breakthrough, a recent follow-up analysis by \cite{Pinto16nat} utilizing the high-resolution reflection grating spectrometer (RGS) aboard \xmm\ was able to resolve this low-energy spectral structure into several discrete emission and absorption features, and found that \ngc\ does indeed exhibit an extreme ionised outflow, potentially consisting of multiple velocity components spanning $\sim$0.2--0.25$c$. Here, by considering the high-energy \xmm\ and \nustar\ data available for \ngc, we report on a detection of an ionized iron \ka\ component to the ultrafast outflow (UFO) discovered by \cite{Pinto16nat}.
\label{sec_dis} We have presented the detection of an absorption feature at $E=8.77^{+0.05}_{-0.06}$ \kev\ in the X-ray spectrum of the ULX \ngc, found by combining data from the \xmm\ and \nustar\ observatories. Owing to the extreme energy of this feature, and the low flux of \ngc, the combination of \xmm\ and \nustar\ is particularly vital to this detection. This provides a broad bandpass, enabling robust continuum estimation both above and below the line energy, and significantly enhances the S/N over what each observatory individually would return in commensurate exposure times. Both of these issues hindered our previous attempt to search for absorption in \ngc\ (using \xmm\ only; \citealt{Walton12ulxFeK}) to the extent that this feature could not be seen. Furthermore, the combination of the different detectors aboard \xmm\ and \nustar\ allows us to effectively rule out an instrumental systematic origin, given that all the detectors utilized show consistent low residuals to the continuum emission. Associating this feature with highly ionised iron, either \fexxv\ or \fexxvi, implies an extreme outflow velocity of 0.2--0.25$c$. Photoionisation modeling marginally prefers a solution in which this absorption is dominated by \fexxv\ ($\log\xi\sim3.3$, $v_{\rm{out}}\sim0.25c$), but with the observational signature of this absorber being dominated by this single line there is significant degeneracy, with solutions dominated by \fexxvi\ ($\log\xi\sim4.5$, $v_{\rm{out}}\sim0.2c$) providing similarly good fits. These velocities are consistent with the UFO recently discovered by \cite{Pinto16nat}, that was identified through the detection of highly blue-shifted absorption lines from moderately ionised material in the low-energy X-ray band, suggesting that we are seeing an iron K$\alpha$ component associated with the same outflow. Ionized iron K$\alpha$ absorption features associated with UFOs ($v_{\rm{out}}>0.1c$) have been seen in several active galaxies (\eg\ \citealt{Tombesi10b}), but never before from an X-ray binary. We note that the energy of the detected feature is just about consistent with the rest-frame energy of the \fexxv\ edge at 8.83\,keV. However, the data prefer the feature to be narrow, and no corresponding absorption line is seen at 6.67\,keV, hence an ionized absorber at rest provides a significantly worse fit in our photoionization modeling ($\Delta\chi^{2}\sim29$). \cite{Pinto16nat} consider two possibilities for the outflow structure: a single zone with low velocity broadening ($v_{\rm{turb}}=20$\,\kmps), and two zones, the second of which has a much higher broadening ($v_{\rm{turb}}=10,000$\,\kmps). The line detected here is narrow; broadening at the latter level seems to be unlikely. We therefore compare our results to the former scenario. Although there is significant degeneracy in our results (Figure \ref{fig_xstar}), the absorption detected here is significantly more ionized, and has a significantly larger column; \citealt{Pinto16nat} found $\log\xi\sim2.3$, and $N_{\rm{H}}\sim2\times10^{22}$\,cm$^{-2}$ for their one-zone model. The absorption detected here may thus arise in a phase of the outflow located closer to the black hole than that contributing the features detected in the RGS. The contrast between $v_{\rm{out}}$ and $v_{\rm{turb}}$ is larger than inferred for the UFOs in PDS456 (\citealt{Nardini15}) and PG1211+143 (\citealt{Pounds03}), despite the similar Gaussian line width constraints. If real, this may provide some clue to the wind geometry, implying that we might not directly view the primary acceleration region, otherwise a smaller contrast would have been expected. However, the constraint on $v_{\rm{turb}}$ is ionization dependent, with the higher ionization solution allowing for $v_{\rm{turb}}$ up to 10,000\,\kmps, more comparable with these other cases. This additional phase of absorption would significantly increase the total mass outflow rate ($\dot{M}_{\rm{out}}$) compared to that inferred from just the low-energy absorption alone. Combining the standard expression for $\dot{M}_{\rm{out}}$ and the definition of the ionisation parameter, we can estimate the kinetic luminosity of the outflow ($L_{\rm{kin}}=1/2\dot{M}v_{\rm{out}}^2$) relative to the bolometric radiative luminosity ($L_{\rm{bol}}$): \begin{equation} \frac{L_{\rm{kin}}}{L_{\rm{bol}}}\approx2{\pi}m_{\rm{p}}\mu\frac{L_{\rm{ion}}}{L_{\rm{bol}}}\frac{v_{\rm{out}}^3}{\xi}{\Omega}C_{\rm{V}} \end{equation} \noindent{where} $m_{\rm{p}}$ is the proton mass, $\mu$ is the mean atomic weight ($\sim$1.2 for solar abundances), $\Omega$ is the (normalized) solid angle subtended by the wind, and $C_{\rm{V}}$ is its volume filling factor (or its `clumpiness'). Although some extrapolation beyond the observed bandpass is necessary, the broadband continuum models constructed by \cite{Bachetti13} and \cite{Miller14} imply $L_{\rm{ion}}/L_{\rm{bol}}\sim0.85$. We therefore find $L_{\rm{kin}}/L_{\rm{bol}}\sim1500{\Omega}C_{\rm{V}}$ and $\sim$60${\Omega}C_{\rm{V}}$ for the lower and higher ionisation solutions, respectively. Thus unless it either has a very small solid angle or a very low volume filling factor (which may be possible; \citealt{King12}), the wind may dominate the energy output from \ngc. While both $\Omega$ and $C_{\rm{V}}$ are unknown, the above $L_{\rm{kin}}/L_{\rm{bol}}$ values are extreme in comparison to similar calculations for even the strongest outflows seen from sub-Eddington systems (\citealt{Blustin05, King12, King14, Nardini15, Miller16}). This is consistent with the basic expectation for a super-Eddington accretion scenario (\citealt{Poutanen07, King09}), as suggested by the unusual broadband X-ray spectrum observed (\citealt{Bachetti13}).
16
7
1607.03124
We present the detection of an absorption feature at E={8.77}<SUB>-0.06</SUB><SUP>+0.05</SUP> keV in the combined X-ray spectrum of the ultraluminous X-ray source NGC 1313 X-1 observed with XMM-Newton and NuSTAR, significant at the 3σ level. If associated with blueshifted ionized iron, the implied outflow velocity is ∼0.2c for Fe xxvi, or ∼0.25c for Fe xxv. These velocities are similar to the ultrafast outflow seen in absorption recently discovered in this source at lower energies by XMM-Newton, and we therefore conclude that this is an iron component to the same outflow. Photoionization modeling marginally prefers the Fe xxv solution, but in either case the outflow properties appear to be extreme, potentially supporting a super-Eddington hypothesis for NGC 1313 X-1.
false
[ "Fe xxv", "Fe xxvi", "lower energies", "the combined X-ray spectrum", "the ultraluminous X-ray source NGC 1313 X-1", "blueshifted ionized iron", "a super-Eddington hypothesis", "the implied outflow velocity", "Eddington", "the outflow properties", "the same outflow", "the Fe xxv solution", "NuSTAR", "absorption", "the ultrafast outflow", "XMM-Newton", "NGC 1313 X-1", "keV", "an iron component", "∼0.25c for Fe xxv" ]
15.748738
8.061106
118
12408446
[ "Liu, Lijuan", "Wang, Yuming", "Wang, Jingxiu", "Shen, Chenglong", "Ye, Pinzhong", "Liu, Rui", "Chen, Jun", "Zhang, Quanhao", "Wang, S." ]
2016ApJ...826..119L
[ "Why is a Flare-rich Active Region CME-poor?" ]
54
[ "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China ; Collaborative Innovation Center of Astronautical Science and Technology, China;", "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China ; Synergetic Innovation Center of Quantum Information &amp; Quantum Physics, University of Science and Technology of China, Hefei, Anhui 230026, China;", "National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, China", "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China ; Synergetic Innovation Center of Quantum Information &amp; Quantum Physics, University of Science and Technology of China, Hefei, Anhui 230026, China", "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China", "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China; Collaborative Innovation Center of Astronautical Science and Technology, China", "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China; Mengcheng National Geophysical Observatory, University of Science and Technology of China, China", "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China; Mengcheng National Geophysical Observatory, University of Science and Technology of China, China", "CAS Key Laboratory of Geospace Environment, Department of Geophysics and Planetary Sciences, University of Science and Technology of China, Hefei, Anhui 230026, China; Collaborative Innovation Center of Astronautical Science and Technology, China" ]
[ "2016AN....337.1082H", "2016ApJ...832..106H", "2017A&A...608A.101P", "2017ApJ...835..211Z", "2017ApJ...843L...9W", "2017ApJ...844..141L", "2017ApJ...849...94J", "2017ApJ...851..133Z", "2017ApJ...851L...1Z", "2017IAUS..327...60C", "2017MNRAS.472..876O", "2017ScChD..60.1383C", "2018AdSpR..61..595K", "2018AdSpR..61.2482S", "2018ApJ...858..121L", "2018ApJ...862...93A", "2018ApJ...867..159S", "2018ApJ...868...59L", "2018ApJ...869..172L", "2018ChJSS..38..665Z", "2018IAUTA..30E...2M", "2018IAUTA..30E...3F", "2018JASTP.174...17Z", "2018JGRA..123.3238W", "2018PhDT.......191H", "2018SoPh..293...16S", "2019ApJ...870...97Z", "2019ApJ...878...38Y", "2019ApJ...881..151L", "2019FrASS...6...44L", "2019JGRA..124.8255L", "2019JKAS...52..133L", "2019MNRAS.486.4936V", "2020AdSpR..65.1641I", "2020ApJ...890...10Z", "2020ApJ...891...54L", "2020ApJ...895...47A", "2020ApJ...900..128L", "2020SCPMA..6369512C", "2020ScChE..63.1699W", "2021ApJ...909..142L", "2021ApJ...917L..29L", "2021JSWSC..11...39G", "2022AN....34310100A", "2022ApJ...926L..14L", "2022ApJ...929....2L", "2022LRSP...19....2C", "2022PhDT........23A", "2022SoPh..297..142C", "2023ApJ...956...24K", "2023FrASS..1035256L", "2023IAUS..372...49L", "2023JASTP.24906106S", "2024ApJ...965L...5K" ]
[ "astronomy" ]
6
[ "Sun: activity", "Sun: coronal mass ejections: CMEs", "Sun: flares", "Sun: magnetic fields", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1969SoPh....6..442S", "1984JFM...147..133B", "1990SoPh..125..219S", "1992ApJ...392..310W", "1994JGR....99.8451F", "1994PhPl....1.1684L", "1995A&A...304..585H", "1995ApJ...440L.109P", "1995SoPh..162..357B", "1996SoPh..168...75A", "1998A&A...339..880Z", "1998ApJ...496L..43B", "1999ApJ...518L..57A", "1999GMS...111.....B", "1999GeoRL..26..627C", "2000ApJ...540..583S", "2000JApA...21..245Z", "2000JGR...105.2375L", "2002ApJ...569.1016F", "2002SoPh..205..325G", "2002SoPh..208...43G", "2002SoPh..209..361T", "2003ApJ...589..644L", "2003ApJ...594.1033N", "2003ApJ...595.1277L", "2003ApJ...595.1296L", "2004ApJ...616L.175N", "2004JGRA..109.7105Y", "2005ApJ...630L..97T", "2005JGRA..11012S05Y", "2006AN....327...36T", "2006ApJ...644..575Z", "2006ApJ...644.1258F", "2006ApJ...644.1273J", "2006ApJ...646L..85Z", "2006ApJ...650L.143Y", "2007A&A...462.1121G", "2007AN....328..743T", "2007AdSpR..39.1467A", "2007AdSpR..39.1674D", "2007ApJ...655L.117S", "2007ApJ...656.1173L", "2007ApJ...661L.109G", "2007ApJ...665.1428W", "2008ApJ...679L.151L", "2008ApJ...680.1516W", "2008ApJ...683.1160Z", "2008SSRv..136....5K", "2008SoPh..248..485O", "2009AdSpR..43..739S", "2010ApJ...708..314A", "2011JGRA..11612108C", "2012A&A...543A..49C", "2012ApJ...759...71E", "2012ApJ...759L...4T", "2012ApJ...761..105L", "2012SoPh..275...17L", "2012SoPh..278..347V", "2014SSRv..186..285P", "2014SoPh..289.3483H", "2014SoPh..289.3549B", "2015ApJ...798..135B", "2015ApJ...801L..23T", "2015ApJ...804L..28S", "2015SoPh..290..811W" ]
[ "10.3847/0004-637X/826/2/119", "10.48550/arXiv.1607.07531" ]
1607
1607.07531_arXiv.txt
\label{sec:intro} Both solar flares and coronal mass ejections (CMEs) indicate the rapid release of a huge amount of magnetic energy in the solar corona, and in particular, CMEs are the most important driving source of hazardous space weather near the geospace. As the major producer of flares and CMEs, active regions (ARs) have been studied for decades. It was revealed based on lots of observational studies that parameters characterizing the AR's non-potentiality, e.g., shear length, magnetic gradient, total electric current, free energy, are all correlated with the flare and CME productivity \citep[e.g.,][]{Canfield_etal_1999, Sammis_etal_2000, Falconer_etal_2002, Falconer_etal_2006,Leka_Barnes_2003a, Leka_Barnes_2003b, Jing_etal_2006, Ternullo_etal_2006, Schrijver_2007, Georgoulis_Rust_2007, Guo_etal_2007, Wang_Zhang_2008}, and particularly larger ARs are more likely to produce eruptions \citep[e.g.,][]{Tian_etal_2002a, Chen_etal_2011}. However, not all large ARs have similar productivities in flares and CMEs, some may be productive in flares only \citep[e.g.,][]{Tian_etal_2002a, Akiyama_etal_2007, ChenA_Wang_2012}. How to distinguish the productivity of an AR is a key issue in space weather forecasting, and still unsolved so far. The recent super AR, 12192, crossing the visible solar disk during 2014 October 17 -- 30, caught a wide attention \citep{RHESSI_Nugget_239, Sun_etal_2015, Thalmann_etal_2015}. It is the largest AR since 1990 November, but produced only one small CME though a total of 127 C-class and intenser flares, including 32 M-class and 6 X-class ones, were generated. Flares and CMEs are thought to be the consequences of the same eruptive process \citep[e.g.,][]{Harrison_1995, Lin_Forbes_2000}. Although the released energy of them during a strong eruption are on the same order of about $10^{32}$ erg \citep{Emslie_etal_2012}, they are clearly different. Flares are relatively local phenomena and the released energy is mostly converted into radiation and energetic particles; while CMEs are more global phenomena and the energy mostly goes into mechanical energies through ejection of magnetized plasma structures. An intense flare may not necessarily be accompanied by a CME \citep[e.g.,][]{Feynman_Hundhausen_1994, Green_etal_2002, Yashiro_etal_2005, Wang_Zhang_2007}, because whether or not there is a CME is substantially determined by the driving force of the inner core magnetic field and the confining force of the external overlying field \citep[e.g.,][]{Wang_Zhang_2007, LiuY_2008, Schrijver_2009}. The inner driver is always in a form of highly sheared or twisted magnetic structure, e.g., a flux rope as required in most CME models \citep{Tamari_etal_1999, Torok_Kliem_2005}. Sheared or twisted field carries magnetic helicity, thus provides a way to transport the helicity naturally \citep{Low_1994, Tamari_etal_1999}. Since magnetic helicity is an invariant in the high conductive corona, it makes a point that a CME may be an inevitable product with the accumulation of helicity in corona \citep{Low_1994, Greenlm_etal_2002, Nindos_etal_2003, Zhang_2006, Zhang_flyer_low_2006, Zhang_Low_2008, Valori_etal_2012, Liu_Schuck_2012b}. However, such an AR of continuously generating M and X-class flares without a strong CME was rarely noticed before, and particularly, 3 out of 6 non-CME X-class flares were of long-duration (lasting more than one hour), which is quite conflicting to many earlier studies that long-duration flares tend to be easier to erupt out \citep[e.g.,][]{Harrison_1995, Yashiro_etal_2006}. In order to understand the underlying physical nature, we compare this super AR with other two pairs of ARs, 11157 and 11158, 11428 and 11429 firstly, then investigate the temporal evolution of photospheric parameters, pre-flare distribution of current helicity, and decay index of the five ARs in the next two sections. In the last section, we give the summary and the discussion.
\label{sec_con} In this work, through comparing AR 12192 with other four ARs, we find that three parameters: the total magnetic flux ($\Phi$), total unsigned vertical current ($I_{total}$), proxy of photospheric free magnetic energy ($\rho_{tot}$), could be responsible for the flare productivity of our sample ARs. The flare-rich only AR 12192, same as the other two flare-rich ARs, 11158 and 11429, has larger $\Phi$, $I_{total}$ and $\rho_{tot}$, which means that they have larger size, and contain stronger current system and more free magnetic energy than the two inert ARs 11157 and 11428. It is reasonable since sufficient amount of free magnetic energy is a necessary condition for an AR to power flares.% No single threshold on any parameter could be used to distinguish the flare and CME productivity of the ARs, but the combination of the mean current helicity and the total unsigned current helicity can be used to distinguish the flare and CME productivity. The magnitude of the mean current helicity ($|\overline{H_c}| $) is large for the CME-rich ARs, and small for AR 12192 and the other two CME-poor ARs, while the total unsigned current helicity (${H_c}_{total}$) of AR 12192 is as large as the two CME-rich ARs, indicating the presence of sheared or twisted field in all three flare-productive ARs. Considering the spatial distribution of current helicity, AR 12192 has $h_c$ concentrated in only one polarity, suggesting the absence of a mature seed structure for CME formation during flares. The CME-rich ARs can also be distinguished by the constraint of the overlying arcade field: AR 12192 has a smaller decay index than the CME-rich ARs, thus no strong CME accompanied the many intense flares it produced. Our study here suggests that pre-existing seed structures at flaring position might be a necessary condition for CMEs. Besides, a large decay index above the AR's flaring neutral lines, which indicates a weak constraint, may be another necessary condition for CMEs. All these facts explain the unusual behaviour of the AR 12192: super flare-rich but CME-poor. The conclusion is obtained based on a sample of five ARs, it's generality should be checked within a larger sample, which would be performed in the future.
16
7
1607.07531
Solar active regions (ARs) are the major sources of two of the most violent solar eruptions, namely flares and coronal mass ejections (CMEs). The largest AR in the past 24 years, NOAA AR 12192, which crossed the visible disk from 2014 October 17 to 30, unusually produced more than one hundred flares, including 32 M-class and 6 X-class ones, but only one small CME. Flares and CMEs are believed to be two phenomena in the same eruptive process. Why is such a flare-rich AR so CME-poor? We compared this AR with other four ARs; two were productive in both and two were inert. The investigation of the photospheric parameters based on the SDO/HMI vector magnetogram reveals that the flare-rich AR 12192, as with the other two productive ARs, has larger magnetic flux, current, and free magnetic energy than the two inert ARs but, in contrast to the two productive ARs, it has no strong, concentrated current helicity along both sides of the flaring neutral line, indicating the absence of a mature magnetic structure consisting of highly sheared or twisted field lines. Furthermore, the decay index above the AR 12192 is relatively low, showing strong constraint. These results suggest that productive ARs are always large and have enough current and free energy to power flares, but whether or not a flare is accompanied by a CME is seemingly related to (1) the presence of a mature sheared or twisted core field serving as the seed of the CME, or (2) a weak enough constraint of the overlying arcades.
false
[ "CME", "productive ARs", "larger magnetic flux", "power flares", "strong constraint", "Flares", "NOAA AR", "coronal mass ejections", "enough current and free energy", "a mature magnetic structure", "AR", "a mature sheared or twisted core field", "CMEs", "highly sheared or twisted field lines", "the overlying arcades", "6 X-class ones", "Solar active regions", "current, and free magnetic energy", "the flaring neutral line", "contrast" ]
12.629894
16.097452
2
12472999
[ "Morselli, L.", "Renzini, A.", "Popesso, P.", "Erfanianfar, G." ]
2016MNRAS.462.2355M
[ "The effect of disc inclination on the main sequence of star-forming galaxies" ]
13
[ "Excellence Cluster Universe, Boltzmannstr. 2, D-85748 Garching bei München, Germany", "INAF-Osservatorio Astronomico di Padova, Vicolo dell'Osservatorio 5, I-35122 Padova, Italy; National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "Excellence Cluster Universe, Boltzmannstr. 2, D-85748 Garching bei München, Germany", "Excellence Cluster Universe, Boltzmannstr. 2, D-85748 Garching bei München, Germany" ]
[ "2017ApJ...848...87C", "2017AstBu..72....1M", "2017MNRAS.468.1850H", "2018A&A...615A...7L", "2018ApJ...856....5Q", "2018ApJ...862....3B", "2018ApJ...862...23Q", "2018MNRAS.480.3788W", "2019MNRAS.489.1265M", "2020MNRAS.498.4345Z", "2021MNRAS.508.4459F", "2022A&A...665A.144M", "2023ApJS..267...17L" ]
[ "astronomy" ]
2
[ "galaxies: bulges", "galaxies: fundamental parameters", "galaxies: spiral", "galaxies: star formation", "Astrophysics - Astrophysics of Galaxies" ]
[ "2000AJ....120.1579Y", "2002ApJS..142....1S", "2004ApJS..153..411D", "2004MNRAS.351.1151B", "2006ApJ...644..792R", "2006ApJ...647..128E", "2007A&A...468...33E", "2007ApJ...660L..43N", "2007ApJ...670..156D", "2007ApJS..173..267S", "2007ApJS..173..293W", "2007ApJS..173..342M", "2009ApJ...698L.116P", "2009ApJS..182..543A", "2010A&A...518L..25R", "2010ApJ...721..193P", "2010MNRAS.404..792M", "2011A&A...532A.145P", "2011ApJ...730...61K", "2011ApJ...739L..40R", "2011ApJ...742...96W", "2011ApJS..196...11S", "2011MNRAS.410..166L", "2012A&A...537A..58P", "2012ApJ...747L..31S", "2012ApJ...754L..14S", "2012ApJ...754L..29W", "2013ApJ...777L...8K", "2013ApJ...778..131L", "2013MNRAS.432..359T", "2013MNRAS.434.2503S", "2013MNRAS.435.2835W", "2014A&A...561A..86M", "2014ApJ...785L..36A", "2014ApJ...791L..25S", "2014ApJ...795..104W", "2014ApJS..214...15S", "2014MNRAS.440..889S", "2014MNRAS.442..509B", "2014MNRAS.443...19R", "2015ApJ...801...80L", "2015ApJ...801L..29R", "2015MNRAS.449..820W", "2015MNRAS.450..435S", "2015MNRAS.450..763M", "2015Sci...348..314T", "2016MNRAS.455.2839E" ]
[ "10.1093/mnras/stw1750", "10.48550/arXiv.1607.05472" ]
1607
1607.05472_arXiv.txt
\label{intro} The stellar mass and star formation rate (SFR) of star-forming (SF) galaxies are tightly correlated with each other and since \cite{Noeske07} such a correlation is designed as the {\it Main Sequence} (MS) of star-forming galaxies. In a series of seminal papers \citep{Noeske07,Daddi07,Elbaz07} it was shown that such tight correlation persists to at least redshift $\sim 2.5$ with nearly constant slope and dispersion. Subsequent studies have confirmed the persistence of a MS relation and extended it all the way to at least $z\sim 4$ and possibly as much as $z\sim 6$ \citep{Pannella09, Peng10, Rodighiero10, Rodighiero11, Rodighiero14, Karim11, Popesso11, Popesso12, Wuyts11, Whitaker12, Whitaker14, Sargent12, Kashino13, Bernhard14, Magnelli14, Speagle14,Steinhardt14}. However, slope, shape, dispersion and redshift evolution of the MS relation can vary substantially from one study to another, with the logarithmic slope of the relation ranging from $\sim 0.4$ up to $\sim 1$, as illustrated by the compilation in \cite{Speagle14}. Much of these differences can be traced back to the criterion used to select galaxies. For example, if galaxies are selected in a SFR-limited fashion, such as in UV- or Far-IR-selected samples, then no MS is recognisable and the SFR remains roughly constant with stellar mass \citep{Erb06,Reddy06, Lee13}, because only galaxies with SFR above threshold are recovered, especially at low stellar masses \citep{ Rodighiero14}. Besides on the first selection of galaxies, the slope of the MS also depends on how galaxies are selected as star forming, e.g., for being bluer than some threshold colour, or for having a specific SFR above a given value, etc. To obviate to some of these limitations, \cite{Renzini15} have proposed to use the {\it ridge line} of the 3D distribution of galaxies in the space having for axes SFR, stellar mass and number of galaxies in SFR-mass bins, or, equivalently, taking the mode of the SFR in each mass bin. With this definition, the resulting MS does not depend on a pre-selection of SF galaxies. Using the database of the Sloan Digital Sky Survey (SDSS) data release 7 (DR7) \citep{Abazajian07}, Renzini \& Peng have derived for the MS of local galaxies a constant slope of $\sim 0.76$, whereas the dispersion may increase somewhat with mass. Using the same database, \cite{Abramson14} have decomposed galaxies in their bulge and disk components, finding that the MS slope gets close to 1 if the SFR is plotted as a function of the {\it disk} mass, as opposed to the total stellar mass. This comes from bulges harbouring very little star formation, if at all, hence contribute mass but not much SFR. Moreover, the fraction of galaxy mass in the bulge component increases steadily with total stellar mass, which then results in a steepening of the MS when the bulge mass is not included. A similar effect (MS slope $\sim 1$) is found for pure disk galaxies at $z\sim1$ \citep{Salmi12}. Deviations of the MS from linearity have also been reported, with either a steepening at low masses (e.g., \citealt{Whitaker14}) or flattening at high masses (e.g., \citealt{Whitaker12,Lee15}). These features may arise from a number of effects, such as the inclusion among SF galaxies of objects in which star formation is being quenched \citep{Renzini15}, or to the development of a passive bulge within the most massive galaxies (e.g., \citealt{Mancini15,Tacchella15, Erfanianfar16}). In this paper we explore the effect of the disk inclination on the MS slope and dispersion. Indeed, reddening must be substantially higher for edge-on galaxies compared to face-on ones, and therefore more uncertain the extinction corrections to apply to the derived SFRs. We also study the contamination of inclined disks in the lower envelope of the MS, hence for the population of galaxies that are most likely characterised by green colours, and are found in the valley between the blue cloud and the red sequence in the colour-magnitude diagram (Schawinski et al. 2014, Lee et al. 2012, Leitner et al. 2012, Smethurst et al. 2015). Here we illustrate the main results of the experiment, having unveiled an inclination effect much stronger than one could have anticipated from the mere extinction correction.
\subsection{Interpretation of the inclination effects} Figure \ref{f6} illustrates the two major inclination effects, which grow more prominent with increasing stellar mass: with increasing disk inclination- (1) the SF peak shifts to lower SFRs, and (2) the strength of the SF peak increases relative to the peak of quenched galaxies. We believe that they reflect two inclination-dependant effects that have a role in determining the total SFR of a galaxy and are dependent on the stellar mass of the galaxy. The first effect is independent of morphology and is related to dust obscuration. For more inclined galaxies as photons have to travel longer distances within the disk itself before reaching the observer, an increasing fraction of the H$\alpha$ and H$\beta$ fluxes gets absorbed within the plane of the galaxy along the line of sight. As discussed in the previous Section, from Fig. \ref{f7} we can distinguish two different ranges of stellar mass where the trends in the position of the $fiber$ SFR peak with increasing disk inclinations are opposite to each other. For low mass galaxies, M$_{\star}<10^{10.25}M_{\odot}$, the peak of the SF gaussian gradually shifts to larger values when going from face-on to edge-on disks. This implies that the shift at lower values of the $total$ SF peak seen for low mass galaxies in Fig. \ref{f6} is actually caused by the aperture correction. At larger stellar masses, the aperture correction acts together with the dust content of the galaxy resulting in a shift of $\sim0.15-0.20$ dex towards lower values of the $total$ SF peak for a galaxy going from a nearly face-on disk to a nearly edge on disk. The second effect is dependant on galaxy morphology and it is more evident for more massive galaxies, since MS massive galaxies are characterised by larger B/T values than MS counterparts at lower stellar mass. In particular, as we discussed earlier, we believe that this effect is due to a trend, with disk inclination, of the bulge and disk fractions that are sampled by the fibers. To show this effect, we make use of the fiber B/T ratio, B/T$_{fiber}$, that is given in the S11 catalogue. As for the total B/T, we use here the B/T$_{fiber}$ computed from the $r-$band. In Fig. \ref{f_ref}, B/T$_{fiber}$ is shown as a function of the disk inclination angle, $i_{S11}$, for galaxies in different bins of total B/T (in different colours) and stellar mass in the range 10$^{10.75}$-10$^{11}M_{\odot}$. As expected, there is a weak trend of decreasing B/T$_{fiber}$ as the disk inclination increases, due to a larger fraction of the disk entering the fiber. This effect is small as the total B/T is fixed and the decrease in B/T$_{fiber}$ with increasing $i_{S11}$ is steeper for galaxies with larger B/T. The slope also varies with stellar mass, and it is steeper for more massive galaxies than for lower mass ones. This is because at fixed B/T the bulge component in more massive galaxies is characterised by a larger radius than the bulge of a lower mass galaxy, hence the variations with disk inclination within the fiber are stronger. In massive, face-on star-forming galaxies the $3''$ fibers sample --in most cases-- just the bulge of the galaxies. Since most bulges are quenched most fibers sample just the quenched part of the galaxies. Apparently, in such cases the aperture correction is not sufficient to properly include the full SFR of the disk. However, as the inclination increases, an increasing fraction of the SF disk enters the fiber, the galaxy is recognised as SF and the aperture correction goes some way towards including the SFR of the disk. Still, quite possibly not all of it, as we know that the aperture correction and dust content are responsible, in massive galaxies, of a shift of $\sim0.15-0.20$ dex towards lower values of the $total$ SF peak. The shift of $\sim 0.4$ dex in the position of the SF peak in the most massive edge-on galaxies indicates that the actual SFR is on average underestimated by a factor $\sim 2.5$ for these galaxies, a quite significant effect. Overall, this extreme inclination effect is relevant for only a small fraction of galaxies, say $\sim 25-30\%$ of the whole sample, and especially confined to the most massive galaxies. Therefore, by and large the bulk of the population is reasonably well accounted for. But minorities count for some interesting issue, as discussed next. \subsection{Effects on the Main sequence} \begin{figure*} \centering \includegraphics[width=0.7\textwidth, keepaspectratio, trim=0cm 0cm 0cm 0cm]{Fig5.eps} \caption{ The SFR-$M_\star$ plane and the MS relation for galaxies that have different disk inclination angles. Each bin has been color-coded as a function of the number of galaxies divided by the average $V_{max}$ value in the bin. The dotted black line is the MS for all galaxies in our final catalogue, i.e. for galaxies with disk inclination angles $0^{\circ}\le i\le 90^{\circ}$. The MS relations are obtained from a linear regression of the mean values of the gaussian fit to the high side of the SFR distributions, as shown in Figure \ref{f5} (red gaussians). The same applies to the dispersion values, shown as vertical bars. The MS fit for the total sample, irrespective of inclination, is shown in black.} \label{f8} \end{figure*} Figure \ref{f8} shows the SFR-$M_{\star}$ plane where each bin in Log(SFR) and Log($M_{\star}$) has been colour coded as a function of the space density in that bin, computed as the ratio of the number of galaxies and the average $V_{max}$, and it is indicated by the greyscale. The best fit MS relations, i.e., Log(SFR)\ =\ $m\, {\rm Log}(M_{\star})+c$, computed in different bins of disk inclination angle, are shown in Figure \ref{f8} with different colours. In each stellar mass bin, the MS value and dispersion are the values written in red in each panel of Figure \ref{f6}, and are the peak and dispersion of the gaussian curve obtained from the fit to the right part of the SF distribution of SFR (gaussians are shown in red in Figure \ref{f6}). The slope and intercept of the MS relation in each disk inclination angle interval are obtained from a linear regression using the MS values in each stellar mass bin, covering the range ($10^{8.5}-10^{11.25}\,\msun$). The resulting SFR-$M_{\star}$ relation for the whole sample of SF galaxies (black dotted line in Figure \ref{f8}) has a slope of $m=0.72\pm0.02$ and an intercept of -7.12. If we limit the stellar mass range to $10^{8.5}-10^{11.0}\,\msun$, hence avoiding the most massive bin, for the whole galaxy sample we obtain $m=0.76\pm0.02$, that is the slope found by \cite{Renzini15}. For the subsample of galaxies characterised by nearly face-on disks ($i<30^{\circ}$; in blue in Figure \ref{f8}) we obtain a slope $m=0.74\pm0.02$, with $c=-7.31$. Galaxies with disk inclination angle between $30^\circ \le i <50^\circ$ (in green) have a slope of $m=0.73\pm0.02$ and $c=-7.17$, while for disk inclination angle between and $50^\circ \le i <70^\circ$ (in yellow) the slope is smaller and the intercept is larger: $m=0.68\pm0.03$, $c=-6.76$. Finally, for nearly edge-on disks, ($i\ge70^{\circ}$; in magenta in Figure \ref{f5}) the slope is flatter, $m=0.56\pm0.03$, and the intercept larger, $c=-5.73$. Hence, the MS becomes progressively flatter and the intercept increases when considering galaxies with increasingly inclined disks. In addition to the change of slope and intercept of the MS for different disk inclination angles, another effect that seems clear from Figure \ref{f8} is that the lower envelope of the MS relation and the green valley region seem to be contaminated by galaxies with highly inclined disks. We analyze this effect in the next Section. \subsection{Contaminating the Green valley} \label{green} The shift to lower SFRs of the SF peak of massive galaxies with increasing inclination has an adverse effect on efforts of identifying really {\it quenching} galaxies, i.e., galaxies that have left the main sequence, their SFR is dropping, and are destined to join the {\it graveyard} of quenched galaxies. Even if the total number of high inclination galaxies is relatively small, within the green valley they may even outnumber the truly quenching galaxies. \begin{figure*} \centering \includegraphics[width=0.75\textwidth, keepaspectratio, trim=0 0cm 0 1cm]{Fig6.eps} \caption{Fraction of galaxies with disk inclination angle $i_{S11}\ge70^{\circ}$ as a function of the distance from the MS of star forming galaxies. In the $left \ panel$ the subsample of galaxies with B/T$\le$0.3 is shown, while the $right\ panel$ shows galaxies with 0.3$<$B/T$\le$0.5. The distance from the MS is computed in bins of 0.2 dex, and it is negative for galaxies above the MS, positive for galaxies below the MS. Different colours mark different stellar mass bins. There is a strong trend of increasing fraction of nearly edge-on galaxies when moving at fixed stellar mass from 0.8 dex above the MS, towards lower SFR, up to 0.8 dex below the MS. This trend is stronger for the lowest B/T galaxies. } \label{f9} \end{figure*} To quantify the contamination of high-inclined disks in the green valley, we plot, for each stellar mass bin, the fraction of galaxies with disk inclination $i\ge70^{\circ}$, $f_{i>70^\circ}$, as a function of the Log(SFR) distance from the MS, $\Delta_{\rm MS}$. Here, the MS values are the ones obtained for the total population of galaxies, regardless of the disk inclination. $\Delta_{\rm MS}$ has been computed in bins that are 0.2 dex wide, and it is negative for galaxies above the MS, positive for galaxies below the MS. Results are shown in Figure \ref{f9}: the left panel shows galaxies with B/T$\le$0.3, while in the right panel galaxies with 0.3$<$B/T$\le$0.5 are shown, to underline trends with morphology. Each stellar mass bin is indicated by a different colour. There is a strong increase of the fraction of nearly edge-on galaxies, $f_{i>70^\circ}$, when moving from the upper envelope of the MS towards its lower envelop. For disk galaxies with M$_{\star}<10^{10.75}$M$_{\odot}$, $f_{i>70^\circ}$ computed 0.8 dex below the MS is 4 times larger than $f_{i>70^\circ}$ at 0.6 dex above the MS. For more massive disk galaxies, the peak value of $f_{i>70^\circ}$ is smaller, $\sim35\%$ for M$_{\star}<10^{10.75-11.0}$M$_{\odot}$ and $\sim27\%$ for M$_{\star}<10^{11.0-11.25}$M$_{\odot}$. In fact, the SFR of the most massive galaxies is derived mainly from the D4000 break as soon as we move below the MS. As shown in Figure 4, for M$_{\star}>10^{10.8}$M$_{\odot}$, the fraction of galaxies that have SFR from $H\alpha$ is $\sim$70$\%$ on the MS, and quickly drops to $\sim$20$\%$ at $\sim$0.8 dex below the MS. Since the SFR from D4000 is less affected by reddening as it is derived from a break of the continuum, the increase of $f_{i>70^\circ}$ from the upper to the lower envelop of the MS is milder for the most massive galaxies. The increase in $f_{i>70^\circ}$ is less steep for galaxies with 0.3$<$B/T$\le$0.5 (right panel in Figure \ref{f9}), where $f_{i>70^\circ}$ computed at 0.8 dex below the MS is a factor from 2 to 3 larger than above MS, and it shows less variation as a function of stellar mass. We believe that the different increase in $f_{i>70^\circ}$ for galaxies with different morphology could be caused by two effects: 1) the disk inclination angle in nearly edge-on disks might be underestimated by S11 for galaxies with $0.4<$B/T$\le$0.5, that mainly populate the lower envelop of the MS, and 2) the increase in the fraction of SFR from D4000 when moving from the MS towards its lower envelop is stronger for intermediate B/T galaxies than for disk-dominated galaxies. It is indeed clear that, at any stellar mass, the fraction of galaxies with highly inclined disks increases when moving from the upper MS envelope toward the lower MS envelope and the green valley region. This rises two important points: 1) there is a fraction highly inclined disks {\it intruders} in the green valley that should be removed if the goal is to identify galaxies in which the quenching is ongoing, and 2) the scatter of the MS seems to be dependent on the disk inclination angle, and hence the relation could be tighter when correcting for the inclination effect. \subsection{Conclusions} We have used the spectroscopic catalogue of SDSS DR7 and the bulge-disk decomposition of Simard et al. (2011) to study the effects of disk inclination on the MS of star-forming galaxies and the contamination of highly inclined disks in the green valley region. The main findings of this work can be summarized as follows: \begin{itemize} \item The peak of the total SFR distribution of star-forming galaxies shifts to lower values with increasing disk inclination angles. Such a shift amounts to $\sim0.2$ dex for Log$(M_{\star}/M_{\odot})<10.25$, and increases to $\sim0.4$ dex for more massive galaxies. The dispersion of the best-fit gaussian to the SF peak slightly increases with increasing disk inclination angles, but this is a small effect (at most $\sim 0.1$ dex in $\sigma$). \item When considering the fiber SFR, in less massive galaxies we observe an opposite trend than the one that characterises the total SFR with increasing disk inclination angle. This implies that, for low stellar masses, the shift of $0.2$ dex towards lower SFR in the total SF peak is caused by the aperture correction applied in Brinchmann et al. (2004). For more massive galaxies, we observe an average shift of the SF peak of 0.2 dex towards lower SFR values when considering nearly edge-on disks as compared to galaxies with 30$^{\circ}<i<$70$^{\circ}$. For galaxies with nearly face-on disks, the peak of the SFR distribution of star forming galaxies is found at lower values than for galaxies with 30$^{\circ}<i<$70$^{\circ}$. This shift to lower fiber SFR values in massive galaxies with nearly face-on disks is due to the presence of a massive non-star-forming bulge in the the SDSS fiber. \item The relative strength of the SF and passive peaks of the SFR distributions is a function of inclination, particularly at high stellar masses, with the SF peak growing more prominent with increasing inclination. This effect reflects a change in the morphology of MS galaxies, that at high stellar masses are characterised by an increase of the bulge component. For a galaxy at a given stellar mass and B/T, the relative fractions of disk and bulge within the fiber are a function on the disk inclination, with the bulge fraction that decreases for increasing disk inclination angle. \item The shift in the SFR peak with increasing inclination is the result of the combination of the different effects explained above. At low stellar masses, MS galaxies are mostly pure disks, and the shift is attributed to the aperture correction that is applied to compute the total SFR from the fiber SFR. At larger stellar masses the dust content of galaxies increases and the final shift of the SF peak in nearly edge on disks is the result of: 1) extinction correction to the SFR being systematically underestimated in high inclination galaxies, 2) aperture correction, and 3) increase of the B/T ratio in MS galaxies. \item For the subsample of nearly edge-on galaxies, the MS relation is significantly flatter than for the total sample, with a slope of $m=0.56\pm0.02$, with respect to the $m=0.74\pm0.02$ for the whole sample. Also the intercept significantly differs, going from -7.31 (whole sample) to -5.37 (nearly edge-on). \item The fraction of galaxies with disk inclination angle $i\ge 70^{\circ}$ is a strong function of the position in the SFR-$M_{\star}$ plane with respect to the MS. In fact, the fraction of disk dominated galaxies with $i\ge 70^{\circ}$ increases by a factor of $\sim 4$ at all stellar masses when moving from 0.8 dex above the MS to 0.8 below it. When considering more bulge dominated galaxies (0.3$<$B/T$\le$0.5), the effect is still important, but the fraction of galaxies with highly inclined disks changes by a factor of $\sim 3$ from the upper to the lower envelope of the MS. \end{itemize} Galaxies characterised by highly inclined disks are likely to be strongly obscured, hence their SFR estimates are more uncertain. \cite{Masters10} and \cite{Sodre13} conclude that significant amounts of dust are present in inclined spirals, and that edge-on disks are the reddest objects in the local Universe. Knowing that, many recent works which focus on the study of quenching mechanism of galaxies based on a colour selection of the galaxy samples tend to remove inclined disks to avoid their obscuration problem (see \cite{Tojeiro13}, \cite{Willett15}). \cite{Schaw14} show that the effect of inclination is small, and that the green valley is not appreciably depopulated after correcting for dust extinction. Still, Schawinski et al. (2014) did not consider green valley galaxies with intermediate morphology, which are the relative majority in the morphological classification of Galaxy Zoo. In principle, the green valley region can provide information on the nature and duration of the process that causes the transition of a galaxy from the blue cloud to the red sequence or, in other terms, from the MS to the passive region of the SFR-$M_{\star}$ plane. For this reason, a careful estimate of the galaxy population of the green valley cannot proceed without taking the disk inclination into account, as we have shown that a large fraction of highly inclined disks tend to fill the green valley and lie in the lower envelope of the MS relation.
16
7
1607.05472
We use the Sloan Digital Sky Survey (York et al.) data base to explore the effect of the disc inclination angle on the derived star formation rate (SFR), hence on the slope and width of the main-sequence (MS) relation for star-forming galaxies. We find that SFRs for nearly edge-on discs are underestimated by factors ranging from ∼0.2 dex for low-mass galaxies up to ∼0.4 dex for high-mass galaxies. This results in a substantially flatter MS relation for high-inclination discs compared to that for less inclined ones, though the global effect over the whole sample of star-forming galaxies is relatively minor, given the small fraction of high-inclination discs. However, we also find that galaxies with high-inclination discs represent a non-negligible fraction of galaxies populating the so-called green valley, with derived SFRs intermediate between the MS and those of quenched, passively evolving galaxies.
false
[ "galaxies", "∼0.4 dex", "high-mass galaxies", "derived SFRs", "star-forming galaxies", "high-inclination discs", "low-mass galaxies", "dex", "MS", "the disc inclination angle", "the derived star formation rate", "a non-negligible fraction", "∼0.4", "SFRs", "quenched, passively evolving galaxies", "al", "SFR", "the small fraction", "factors", "the global effect" ]
12.419028
7.647706
-1
12515467
[ "Chan, Man-Ho" ]
2016RAA....16..163C
[ "Does the gamma-ray signal from the central Milky Way indicate Sommerfeld enhancement of dark matter annihilation?" ]
0
[ "Department of Science and Environmental Studies, The Education University of Hong Kong, Tai Po, New Territories, Hong Kong, China" ]
null
[ "astronomy" ]
3
[ "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1931AnP...403..257S", "1997ApJ...490..493N", "2001MNRAS.320L...1S", "2003MNRAS.340..657D", "2005A&A...435....1P", "2005AJ....129.2119S", "2006JETPL..84...45V", "2010PhRvD..81h3502Z", "2011AJ....141..193O", "2011PhRvD..83h3507C", "2011PhRvL.106q1302L", "2012JCAP...07..017A", "2012MNRAS.422.1231G", "2013JCAP...07..016N", "2013PDU.....2..118H", "2013PhRvD..88h3009H", "2013PhRvD..88h3521G", "2013arXiv1307.6862H", "2014MNRAS.437..415D", "2014Natur.506..171P", "2014PhRvD..89d2001A", "2014PhRvD..89f3530Y", "2014PhRvD..90b3526A", "2014PhRvD..90e5002I", "2015ApJ...808..158B", "2015JCAP...12..001P", "2015MNRAS.448L..87C", "2015MNRAS.453..849B", "2015NatPh..11..245I", "2015PhRvD..91f3003C", "2015PhRvL.115w1301A", "2016ApJ...819...44A", "2016PDU....12....1D", "2016PhRvD..93j3512E" ]
[ "10.1088/1674-4527/16/10/163", "10.48550/arXiv.1607.02246" ]
1607
1607.02246_arXiv.txt
% \label{sect:intro} In the past few years, some excess GeV gamma rays emitted from our Galactic center were reported (\cite{Hooper,Huang,Fermi}). The large diffuse signal of GeV gamma rays is hard to be explained by the cosmic ray and pulsar emission. Recent studies point out that the millisecond pulsars can only account no more than 10\% of the GeV excess (\cite{Hooper2,Daylan}). Therefore, the possibility of the emission of gamma rays due to the dark matter annihilation has become a hot topic in the recent years (\cite{Daylan,Gordon,Abazajian,Izaguirre,Calore}). In particular, \cite{Daylan,Calore} discover that the gamma-ray spectrum obtained from Fermi-LAT can be well fitted with $m=30-70$ GeV dark matter annihilation through $b\bar{b}$ channel. The cross section obtained $<\sigma v>=(1.4-2.0) \times 10^{-26}$ cm$^3$ s$^{-1}$ generally agrees with the expected canonical thermal relic abundance cross section $<\sigma v> \approx (2-3) \times 10^{-26}$ cm$^3$ s$^{-1}$. Moreover, the inner slope of the radial-dependence of the gamma-ray emission is $\gamma \approx 1.1-1.3$ (the best-fit value is $\gamma=1.26$), which is consistent with the theoretical expectation from numerical simulations ($\gamma=1-1.5$) (\cite{Daylan}). This work is further supported by a later study which includes the consideration of foreground and background uncertainties (\cite{Calore}). On the other hand, \cite{Chan} shows that this dark matter model can also explain the origin of hot gas near the Galactic center. Therefore, this dark matter model becomes one of the most popular models in dark matter astrophysics. Besides the detection of gamma-ray emission from Galactic center, Fermi-LAT also obtains some upper limits of gamma-ray emission from dwarf galaxies and galaxy clusters. If we assume that the gamma-ray emission is due to the annihilation of $m=40$ GeV dark matter with $b\bar{b}$ channel, the corresponding upper limits of cross sections are $<\sigma v> \approx 1 \times 10^{-26}$ cm$^3$ s$^{-1}$ (\cite{Ackermann,Ackermann2}) and $<\sigma v> \approx (2-3) \times 10^{-25}$ cm$^3$ s$^{-1}$ (\cite{Ando}) for dwarf galaxies and galaxy clusters respectively. In fact, the results obtained in \cite{Daylan,Calore} assume that the annihilation cross section is constant (velocity-independent). However, it has been suggested that the annihilation cross section can be velocity-dependent. For example, the multiple exchange of some light force-carrier particle between the annihilating dark matter particle (the Sommerfeld enhancement) gives $<\sigma v> \propto v^{-\alpha}$, where $\alpha=1$ and $\alpha=2$ for non-resonance and resonance respectively (\cite{Sommerfeld,Zavala,Yang}). Furthermore, the inner slope of dark matter in our Galactic center revealed in \cite{Daylan,Calore} (best-fit $\gamma=1.26$) is a bit too large, compared with the recent observations in Milky Way (\cite{Pato}). Recent studies point out that the Milky Way dark matter density is well-fitted by a NFW density profile ($\gamma=1$) (\cite{Navarro,Iocco,Pato}). Detailed analyses in \cite{Pato} show that the best-fit $2\sigma$ range of the inner slope for the most representative baryonic model is $\gamma=0.6-0.8$. If we assume a generalized NFW profile with local density $\rho_{\odot}=0.4$ GeV cm$^{-3}$, $\gamma>1.2$ is excluded (outside the $5\sigma$ region) for this representative baryonic model. Although these results do not really rule out the possibility of having $\gamma=1.1-1.3$ (some baryonic models can still generate these values), such a large inner slope in Milky Way is certainly questionable. In fact, most of the inner slopes of dark matter density profiles observed do not show $\gamma>1$. For example, most galaxy clusters give $\gamma \approx 1$ (\cite{Pointecouteau}) and most galaxies and dwarf galaxies give $\gamma \le 1$ (\cite{Salucci2,Oh,Loeb}). Furthermore, recent numerical simulations show that baryonic feedback can decrease the inner slope of dark matter such that $\gamma<1$ for normal galaxies (\cite{Governato,Pontzen}). Therefore, the inner slope obtained in \cite{Daylan,Calore} does not give a good agreement with many other observations and recent numerical simulations. In this article, we show that the result obtained in \cite{Daylan} is completely compatible with a velocity-dependent annihilation cross section. If we assume that the dark matter annihilation in the Milky Way center is Sommerfeld-enhanced, the resulting inner slope $\gamma$ obtained would give $\gamma=0.85-1.05$, which agrees with the standard NFW profile ($\gamma=1$). Also, we show that our model satisfies the Fermi-LAT results of nearby dwarf galaxies.
Previously, \cite{Daylan} show that the GeV gamma-ray excess can be explained by the annihilation of $\sim 40$ GeV dark matter through $b\bar{b}$ channel. Based on the morphology of the gamma-ray flux, the best-fit inner slope of dark matter density profile is $\gamma=1.26$. However, recent analyses show that the best-fit $2\sigma$ range of inner slope for most the representative baryonic model is $\gamma=0.6-0.8$ (\cite{Pato}). Also, many observations indicate $\sigma \le 1$ (\cite{Salucci2,Oh,Loeb}). In this article, we show that the GeV gamma-ray excess can also be explained by the Sommerfeld-enhanced dark matter annihilation through $b\bar{b}$ channel with $\gamma=1$ (the NFW profile). In general, our model is compatible with the range of the inner slope $\gamma=0.85-1.05$. By using the results in \cite{Daylan}, we also constrain the parameters of the Sommerfeld enhancement: $\alpha=1$ and $<\sigma v>_0v_0 \approx (2.2-3.2) \times 10^{-19}$ cm$^4$ s$^{-1}$ for $\gamma=1$. Although the annihilation model with Sommerfeld enhanced cross section is more complicated, this model can fully explain the morphology of the gamma-ray flux and favor the smaller inner slope of the dark matter density in Milky Way. Since the morphology of the gamma-ray flux gives $F \propto r^{-(2.2-2.6)}$ (\cite{Daylan}), the inner slope obtained from this model is $\gamma=0.85-1.05$, which gives a very good agreement with the recent analysis in \cite{Pato}. However, if we assume the constant cross section for dark matter annihilation, the required inner slope is $\gamma=1.1-1.3$, which does not satisfy with the observed $2\sigma$ range of the inner slope $\gamma=0.6-0.8$ for the most representative baryonic model (\cite{Pato}). Therefore, our model can alleviate the tension between the existing dark matter annihilation model and the observations. In fact, the Sommerfeld enhancement would greatly enhance the dark matter annihilation rate near the dwarf galactic center because the velocity dispersion is very small there. Therefore, we predict that a strong signal of gamma-ray flux at the dwarf galactic center would be resulted if our model is correct. We show that the gamma-ray fluxes emitted due to the Sommerfeld-enhanced dark matter annihilation from the dwarf galaxies generally satisfy the current upper limits obtained by the 6-year Fermi-LAT data. If the Fermi-LAT can further constrain the upper limits of the gamma-ray flux or the detected gamma-ray spectrum in the future, we can get a tighter constraint on the annihilation cross section as well as the dark matter rest mass.
16
7
1607.02246
Recently, some studies showed that the GeV gamma-ray excess signal from the central Milky Way can be explained by the annihilation of ∼ 40 GeV dark matter through the bb¯ channel. Based on the morphology of the gamma-ray flux, the best-fit inner slope of the dark matter density profile is γ = 1.26. However, recent analyses of the Milky Way dark matter profile favor γ = 0.6 - 0.8. In this article, we show that the GeV gamma-ray excess can also be explained by the Sommerfeld-enhanced dark matter annihilation through the bb¯ channel with γ = 0.85 - 1.05. We constrain the parameters of the Sommerfeld-enhanced annihilation by using data from Fermi-LAT. We also show that the predicted gamma-ray fluxes emitted from dwarf galaxies generally satisfy recent upper limits on gamma-ray fluxes detected by Fermi-LAT.
false
[ "the Milky Way dark matter profile", "recent upper limits", "the dark matter density profile", "Milky Way", "dwarf galaxies", "the Sommerfeld-enhanced dark matter annihilation", "gamma-ray fluxes", "recent analyses", "Fermi-LAT", "Sommerfeld", "the bb¯ channel", "data", "the GeV gamma-ray excess signal", "the predicted gamma-ray fluxes", "the central Milky Way", "the gamma-ray flux", "the GeV gamma-ray excess", "the Sommerfeld-enhanced annihilation", "∼ 40 GeV dark matter", "the annihilation" ]
7.850626
-1.30261
53
1749239
[ "Bambi, Cosimo", "Cárdenas-Avendaño, Alejandro", "Dauser, Thomas", "García, Javier A.", "Nampalliwar, Sourabh" ]
2017ApJ...842...76B
[ "Testing the Kerr Black Hole Hypothesis Using X-Ray Reflection Spectroscopy" ]
133
[ "Center for Field Theory and Particle Physics and Department of Physics, Fudan University, 200433 Shanghai, China ; Theoretical Astrophysics, Eberhard-Karls Universität Tübingen, D-72076 Tübingen, Germany;", "Programa de Matemática, Fundación Universitaria Konrad Lorenz, 110231 Bogotá, Colombia", "Remeis Observatory &amp; ECAP, Universität Erlangen-Nürnberg, D-96049 Bamberg, Germany", "Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA", "Center for Field Theory and Particle Physics and Department of Physics, Fudan University, 200433 Shanghai, China" ]
[ "2016JCAP...09..014N", "2016JCAP...10..003C", "2016PhRvD..94j4062G", "2017ApJS..229...40P", "2017CQGra..34k5003A", "2017EPJC...77..461I", "2017JCAP...08..014S", "2017JPhCS.942a2004B", "2017JPhCS.942a2017T", "2017PhRvD..95f4006B", "2017PhRvD..95j4035Z", "2017PhRvD..95j4043Z", "2018ApJ...865..134X", "2018CQGra..35r5014P", "2018EPJC...78..376Z", "2018GReGr..50..100K", "2018JCAP...08..044L", "2018PhLB..781..626N", "2018PhRvD..97b4043G", "2018PhRvD..97l4005T", "2018PhRvD..98b3018T", "2018PhRvD..98b4007Z", "2018PhRvD..98d4024Y", "2018PhRvL.120e1101C", "2018Univ....4...79B", "2018rnls.confE..37T", "2019ApJ...874..135T", "2019ApJ...875...41Z", "2019ApJ...875...56T", "2019ApJ...878...91A", "2019ApJ...879...80C", "2019ApJ...884..147Z", "2019EL....12530002Z", "2019PhRvD..99h3001T", "2019PhRvD..99j4009N", "2019PhRvD..99j4031Z", "2019PhRvD..99l3007L", "2019PhRvD..99l4026S", "2019PhRvD.100b4039C", "2019PhRvD.100d4026V", "2019PhRvD.100h4055X", "2019arXiv190508012A", "2019rpra.conf....7A", "2020ApJ...895...61R", "2020ApJ...896..160L", "2020ApJ...897...84T", "2020ApJ...899...80A", "2020CQGra..37m5008C", "2020EPJC...80..400Z", "2020EPJC...80..622Z", "2020JCAP...05..026W", "2020JCAP...07..058G", "2020MNRAS.491..417R", "2020MNRAS.496..497Z", "2020MNRAS.498.3565T", "2020PhRvD.101d3010Z", "2020PhRvD.101f4030T", "2020PhRvD.101l3014C", "2020PhRvD.101l4024B", "2020PhRvD.102d4013N", "2020PhRvD.102j3009T", "2020PhRvD.102j4035N", "2020PhRvD.102l4071N", "2020arXiv201207469R", "2020mbhe.confE..28B", "2021ARep...65..902B", "2021ApJ...907...31T", "2021ApJ...910...49R", "2021ApJ...910...52C", "2021ApJ...913...79T", "2021ApJ...913..129T", "2021ApJ...923..175A", "2021EPJC...81..269N", "2021Galax...9...63N", "2021Galax...9...75R", "2021JCAP...01..047Z", "2021JCAP...07..002T", "2021JCAP...09..028K", "2021MNRAS.500..481B", "2021MNRAS.506.4960H", "2021PhRvD.103b4055Z", "2021PhRvD.103j3023A", "2021PhRvD.104b4058A", "2021PhRvD.104d4001R", "2021PhRvD.104h4035Y", "2021SSRv..217...65B", "2021Univ....7..136B", "2021arXiv210311365B", "2021arXiv210604084B", "2021arXiv210801190N", "2021arXiv210803573N", "2021bhns.confE...1K", "2022ApJ...924...72Z", "2022ApJ...925...51R", "2022ApJ...938...53S", "2022EPJC...82..708G", "2022JCAP...01..019T", "2022JCAP...10..040R", "2022MNRAS.512.2082L", "2022MNRAS.517.5721M", "2022PhRvD.105f4073B", "2022PhRvD.105h4046V", "2022PhRvD.105j4004S", "2022PhRvD.106f3009N", "2022PhRvD.106h4041L", "2022Univ....8..626C", "2022arXiv220610474N", "2022arXiv221005322B", "2022arXiv221011959B", "2022hxga.book...81A", "2023ApJ...945...65Y", "2023ApJ...951...19L", "2023CQGra..40t5011B", "2023EPJC...83..321H", "2023EPJC...83..838R", "2023JCAP...08..018C", "2023MNRAS.521..708H", "2023PhRvD.107l4063F", "2023PhRvD.108d4078R", "2023PhRvD.108h3036T", "2023PhRvD.108h4055D", "2023Symm...15.1277B", "2023arXiv230609673R", "2023arXiv231205857B", "2023arXiv231211430E", "2024ApJ...967...35L", "2024EPJC...84..302S", "2024MNRAS.528.2015M", "2024PhRvD.109f4052C", "2024PhRvD.109f4059Z", "2024arXiv240108545T", "2024arXiv240601226M", "2024mbhe.confE..16B" ]
[ "astronomy" ]
9
[ "accretion", "accretion disks", "black hole physics", "gravitation", "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1920RSPTA.220..291D", "1971PhRvL..26..331C", "1972PhRvD...5.2419P", "1974ApJ...191..499P", "1975ApJ...202..788C", "1975PhRvL..34..905R", "1989MNRAS.238..729F", "1993PhRvD..47.5259M", "1995CoPhC..88..109S", "1997ApJ...482L.155Z", "2001ApJS..133..221K", "2002NuPhB.626..377T", "2003IJMPD..12...63L", "2004ApJS..155..675K", "2005ApJS..158...68G", "2006ApJ...636L.113S", "2006ApJ...652.1028B", "2008PhRvL.100k9902P", "2009ApJS..185..477G", "2009GReGr..41.1795S", "2009JCAP...09..013B", "2009JInst...4.4001L", "2010ApJ...718..446J", "2010ApJ...718..695G", "2010MNRAS.409.1534D", "2010SSRv..157..167F", "2011ApJ...731..121B", "2011MNRAS.415.2942C", "2012ApJ...761..174B", "2012LRR....15....7C", "2013ApJ...768..146G", "2013ApJ...773...57J", "2013ForPh..61..742D", "2013JCAP...08..055B", "2013LRR....16....9Y", "2013MNRAS.430.1694D", "2013MNRAS.434.1955D", "2013PhLB..719..419D", "2013PhRvD..87b3007B", "2013PhRvD..88d4002J", "2013PhRvD..88f4022B", "2014ApJ...782...76G", "2014ApJ...797...78K", "2014LRR....17....4W", "2014PhRvD..89j4059B", "2014PhRvD..89l7302B", "2014PhRvD..90f4002J", "2014PhRvD..90l4033G", "2014PhRvL.112v1101H", "2014SSRv..183..277R", "2014SSRv..183..295M", "2015ApJ...811..130J", "2015JCAP...05..025J", "2016CQGra..33e4001Y", "2016CQGra..33f4001B", "2016CQGra..33l4001J", "2016EL....11630006B", "2016JCAP...07..049N", "2016PhLB..756..350K", "2016PhLB..760..254C", "2016PhRvD..93f4015K", "2016PhRvD..93l3008J", "2016SPIE.9905E..1QZ", "2017RvMP...89b5001B" ]
[ "10.3847/1538-4357/aa74c0", "10.48550/arXiv.1607.00596" ]
1607
1607.00596_arXiv.txt
} The theory of general relativity was proposed by Einstein about a century ago and is still the standard framework for the description of the gravitational field and the chrono-geometrical structure of the spacetime. The first test of general relativity can be dated back to the measurement of light bending by the Sun by Eddington in 1919~\citep{edd1919}. Especially over the past 60~years, there have been significant efforts to test the theory in weak gravitational fields, mainly with precise experiments in the Solar System and accurate radio observations of binary pulsars~\citep{will}. Tests of general relativity in the strong gravity regime are nowadays the new frontier, both with electromagnetic radiation~\citep{er1,er2,j2016} and gravitational waves~\citep{gw1,gw2}. Astrophysical black holes are the ideal laboratory for testing strong gravity. In 4-dimensional general relativity, an uncharged black hole is described by the Kerr solution\footnote{There are a number of assumptions behind this statement. In particular, the spacetime must have 4~dimensions and be stationary and asymptotically flat; the exterior must be regular (no singularities or closed time-like curves); the metric is a vacuum solution of the Einstein equations. See, e.g., \citet{nh-rev} for more details.} and is completely described by only two parameters, namely the mass $M$ and the spin angular momentum $J$ of the object. This is the result of the ``no-hair theorem''~\citep{nh1,nh2}. It is remarkable that the spacetime around astrophysical black holes should be well described by the Kerr metric. As soon as a black hole is formed, initial deviations from the Kerr solution are quickly radiated away with the emission of gravitational waves~\citep{k1}. The equilibrium electric charge is extremely small for macroscopic objects and completely negligible for the spacetime geometry~\citep{k2}. Accretion disks typically have a mass of several orders of magnitude smaller than the central object and their impact on the background metric can be safely ignored~\citep{k3,k4}. Within Einstein's theory of gravity, the Kerr metric should well describe the spacetime around astrophysical black holes. Nevertheless, macroscopic deviations from the Kerr spacetime are possible in many scenarios. For instance, \citet{carlos} have recently discovered a family of hairy black holes in 4-dimensional Einstein gravity minimally coupled to a complex, massive scalar field. Hairy black holes generically arise when scalar fields non-minimally coupled to gravity, and an example is the dilaton in Einstein-dilaton-Gauss-Bonnet gravity~\citep{hbh-EdGB}. Quantum gravity effects might also produce macroscopic corrections to the Kerr metric~\citep{dvali1,dvali2,giddings}. Electromagnetic and gravitational radiations can test general relativity in different ways. The properties of the electromagnetic radiation emitted by the accreting gas close to a black hole depend on both the gas motion in the strong gravity region and the photon propagation from the emission point in the disk to the detection point in the flat faraway region. In this case, we can test the Kerr metric as Solar System experiments have so far tested the Schwarzschild solution in the weak field limit. However, it is not possible to distinguish a Kerr black hole in general relativity from a Kerr black hole in an alternative metric theory of gravity, because the geodesic motion is the same~\citep{psaltis}. Gravitational waves can instead probe the field equations of the theory, while they are less suitable to perform model-independent tests. The two approaches can thus be seen as complementary; see, e.g., \citet{kz}, \citet{comp1}, and \citet{comp2}. With the electromagnetic approach, currently there are two leading techniques to probe the strong gravity region around a black hole: the study of the thermal spectrum of thin disks (continuum-fitting method)~\citep{cfm1,cfm1b,cfm2} and the analysis of the relativistically smeared reflection spectrum of thin disks (reflection method)~\citep{iron1,iron1b,iron2}. Both techniques have been developed for measuring black hole spins under the assumption of Kerr background and can be naturally extended for testing the Kerr metric~\citep{cfm3,cfm4,cfm5,iron3,ss09,iron4,iron5,cfm-iron-c,iron6,yy16}. The reflection method has a number of advantages with respect to the study of the thermal spectrum. It can be easily applied to both stellar-mass and supermassive black holes\footnote{The continuum-fitting method has also been applied to supermassive black holes, but only in very special cases~\citep[e.g.][]{czerny2011,done2013}.}. It is independent of the black hole mass and distance, while the inclination angle of the disk with respect to the line of sight of the observer can be inferred from the fit of the reflection spectrum; with the continuum-fitting method, these three quantities have to be obtained from other measurements and their uncertainty is often large. In the presence of high quality data and the correct astrophysical model, the reflection method is potentially quite a powerful tool to constrain the metric around black holes~\citep[see, for instance,][]{jjc1,jjc2,jjc3,comp1}. Theoretical models of X-ray reflection have been undergoing active development over three decades~\citep[see][for a review]{fr2010}. Currently, the most advanced model is {\sc xillver}~\citep{xillver1,xillver2}, and its relativistic counterpart {\sc relxill}~\citep{relline2,relxill}. These are the state-of-the-art in modeling reflection in strong gravity. Compared to all earlier reflection codes, {\sc xillver} provides a superior treatment of the radiative transfer, as well as an improved calculation of the ionization balance, by implementing the photoionization routines from the {\sc xstar} code~\citep{kallman2001}, which incorporates the most complete atomic database for modeling synthetic photoionized X-ray spectra. The microphysics captured by {\sc xillver} is much more rigorous than for any earlier code, principally because of the detailed treatment of the K-shell atomic properties of the prominent ions~\citep[e.g.,][]{garcia2005,kallman2004,garcia2009}. The model {\sc relxill} is the result from the combination of {\sc xillver} with the relativistic blurring code {\sc relconv}~\citep{relline1}. {\sc relconv} is a relativistic convolution code that, assuming the Kerr metric, requires as input the local spectrum at any emission point in the disk and gives as output the spectrum measured by a distant observer. The aim of our work here is to construct a model to extend {\sc relxill} to a generic stationary, axisymmetric and asymptotically flat black hole metric. We replace {\sc relconv} with a more general relativistic convolution code, while we maintain {\sc xillver} because the microphysics of the local spectrum does not change. In this Paper, we present a new code to compute transfer functions in any stationary, axisymmetric, and asymptotically flat black hole metric and extend {\sc relxill} for testing the Kerr black hole hypothesis. Current studies along this line of research model the X-ray spectrum with a simple power-law plus a relativistically broadened iron line~\citep{jjc1,jjc2,jjc3}. This can be sufficient for a preliminary study and a qualitative analysis. However, this is definitively not adequate if we really want to test general relativity. Here we employ the formalism of the transfer function for thin accretion disks~\citep{cun75}. In this framework, the calculations of the reflection spectrum are split into two parts: the calculation of the transfer function and the calculation of the reflection spectrum in the rest-frame of the gas. The transfer function only depends on the metric of the background and takes into account all the relativistic effects (gravitational redshift, Doppler boosting, light bending). The local spectrum is obtained by solving radiation transfer on a plane-parallel, 1-dimensional slab and is not strictly related to the metric of the spacetime. In order to test the Kerr metric, our model must be able to compute the X-ray reflection spectrum of a thin disk in a background more general than the Kerr solution and that includes the Kerr solution as a special case. The test-metric is described by the mass $M$ and the spin angular momentum $J$ of the object, as well as by a number of ``deformation parameters''. The latter are used to quantify possible deviations from the Kerr metric and are the parameters to constrain from observations to verify the Kerr black hole hypothesis. The Kerr metric is recovered when all the deformation parameters vanish, while there are deviations from the Kerr solution in the presence of at least one non-vanishing deformation parameter. In the standard case of the Kerr metric, the calculations of the transfer function exploit some specific properties of the Kerr solution~\citep{cun75,spe95}. Because of the presence of the Carter constant, the equations of motions are separable. More importantly, the equations in the $(r,\theta)$ plane can be reduced to elliptic integrals. This significantly simplifies the calculations of the transfer function. In our more general case, the transfer function is evaluated by integrating the photon geodesic equations from the point of detection in the plane of the distant observer backward in time to the point of emission in the accretion disk. Our calculations are inevitably longer than those in the Kerr metric that solve elliptic integrals. The Paper is organized as follows. In Section~\ref{s-2}, we review the formalism of the transfer function and, in Section~\ref{s-6}, the Johannsen metric~\citep{j-m}, which is the one adopted in our current version of the code. Section~\ref{s-4} describes our numerical method to compute the transfer function. In Section~\ref{s-5}, we compare transfer functions and single iron line shapes produced by our code for a few Kerr solutions with those calculated by existing codes. Section~\ref{s-5bis} shows some examples of transfer functions and single line shapes in the Johannsen metric. In Section~\ref{s-new}, we simulate several observations of a bright black hole binary with NuSTAR and LAD/eXTP and we fit the data with our new version of {\sc relxill} to constrain one of the deformation parameters in the Johannsen metric as an illustrative example of the application of the new model and the constraining power of current and future X-ray missions. Summary and conclusions are reported in Section~\ref{s-7}. In Appendix~\ref{s-3}, we present all the formulas to compute the transfer function for a thin accretion disk in a generic stationary, axisymmetric, and asymptotically flat black hole spacetime. Throughout the Paper, we employ units in which $G_{\rm N} = c = 1$ and the convention of a metric with signature $(-+++)$. Except in Section~\ref{s-6}, we set the black hole mass parameter $M$ as defined in the Kerr and Johannsen metrics equal to 1.
\label{s-7} In this Paper, we present the first X-ray reflection model for testing the spacetime metric around astrophysical black holes. Previous work suggests that the reflection method is quite a promising technique to test the Kerr black hole hypothesis with electromagnetic radiation. However, current studies employ simplified models. In the best cases, the X-ray spectrum is approximated by a power law with an iron line. Similar models can work for preliminary studies, but they are definitively inadequate to perform precise tests of general relativity in the strong gravity regime. {\sc relxill} is currently the most sophisticated model to fit the X-ray reflection spectrum of black holes under the assumption that the spacetime is described by the Kerr solution. It results from the combination of the relativistic convolution model for the Kerr metric {\sc relconv} and the reflection code for the local spectrum {\sc xillver}. By calculating the transfer function for a generic background, we have a new relativistic convolution model to replace {\sc relconv}. After merging our new relativistic convolution model with {\sc xillver}, we obtain the extension of {\sc relxill} to generic stationary, axisymmetric, and asymptotically flat black hole spacetime. We have described our new code and the relevant formulas for the calculation of the transfer function. We have shown that our calculations reach the necessary accuracy for our tests. We have simulated some observations of a bright black hole binary with NuSTAR and LAD/eXTP to illustrate the constraining power of current and future X-ray missions. The current version of the code adopts the Johannsen metric, but it is straightforward to employ any other stationary, axisymmetric, and asymptotically flat black hole metric. Work on other non-Kerr metrics, such as that proposed in~\citet{ref-krz}, is currently underway. In a forthcoming paper, we will apply our new model to a specific source for constraining the deformation parameters of the Johannsen metric from available X-ray data.
16
7
1607.00596
We present the first X-ray reflection model for testing the assumption that the metric of astrophysical black holes is described by the Kerr solution. We employ the formalism of the transfer function proposed by Cunningham. The calculations of the reflection spectrum of a thin accretion disk are split into two parts: the calculation of the transfer function and the calculation of the local spectrum at any emission point in the disk. The transfer function only depends on the background metric and takes into account all the relativistic effects (gravitational redshift, Doppler boosting, and light bending). Our code computes the transfer function for a spacetime described by the Johannsen metric and can easily be extended to any stationary, axisymmetric, and asymptotically flat spacetime. Transfer functions and single line shapes in the Kerr metric are compared to those calculated from existing codes to check that we reach the necessary accuracy. We also simulate some observations with NuSTAR and LAD/eXTP and fit the data with our new model to show the potential capabilities of current and future observations to constrain possible deviations from the Kerr metric.
false
[ "light bending", "Transfer functions", "Kerr", "astrophysical black holes", "existing codes", "the Kerr metric", "possible deviations", "gravitational redshift", "Johannsen", "the Johannsen metric", "single line shapes", "the Kerr solution", "Doppler", "the metric", "the first X-ray reflection model", "the necessary accuracy", "Cunningham", "a thin accretion disk", "current and future observations", "The transfer function" ]
8.849474
2.804938
61
4017664
[ "Hinkel, Natalie R.", "Young, Patrick A.", "Pagano, Michael D.", "Desch, Steven J.", "Anbar, Ariel D.", "Adibekyan, Vardan", "Blanco-Cuaresma, Sergi", "Carlberg, Joleen K.", "Delgado Mena, Elisa", "Liu, Fan", "Nordlander, Thomas", "Sousa, Sergio G.", "Korn, Andreas", "Gruyters, Pieter", "Heiter, Ulrike", "Jofré, Paula", "Santos, Nuno C.", "Soubiran, Caroline" ]
2016ApJS..226....4H
[ "A Comparison of Stellar Elemental Abundance Techniques and Measurements" ]
69
[ "School of Earth &amp; Space Exploration, Arizona State University, Tempe, AZ 85287, USA", "School of Earth &amp; Space Exploration, Arizona State University, Tempe, AZ 85287, USA", "School of Earth &amp; Space Exploration, Arizona State University, Tempe, AZ 85287, USA", "School of Earth &amp; Space Exploration, Arizona State University, Tempe, AZ 85287, USA", "School of Earth &amp; Space Exploration, Arizona State University, Tempe, AZ 85287, USA", "Instituto de Astrofísica e Ciências do Espaço, Universidade do Porto, CAUP, Rua das Estrelas, 4150-762 Porto, Portugal", "Observatoire de Genève, Université de Genève, CH-1290 Versoix, Switzerland", "NASA Goddard Space Flight Center, Code 667, Greenbelt MD 20771, USA ; Department of Terrestrial Magnetism, Carnegie Institution of Washington, 5241 Broad Branch Road, NW, Washington DC 20015, USA", "Instituto de Astrofísica e Ciências do Espaço, Universidade do Porto, CAUP, Rua das Estrelas, 4150-762 Porto, Portugal", "Research School of Astronomy &amp; Astrophysics, Australian National University, Cotter Road, Weston Creek, ACT 2611, Australia", "Department of Physics and Astronomy, Uppsala University, Box 516, 75120 Uppsala, Sweden", "Instituto de Astrofísica e Ciências do Espaço, Universidade do Porto, CAUP, Rua das Estrelas, 4150-762 Porto, Portugal", "Department of Physics and Astronomy, Uppsala University, Box 516, 75120 Uppsala, Sweden", "Department of Physics and Astronomy, Uppsala University, Box 516, 75120 Uppsala, Sweden; Lund Observatory, Department of Astronomy and Theoretical Physics, Box 43, 221 00, Lund, Sweden", "Department of Physics and Astronomy, Uppsala University, Box 516, 75120 Uppsala, Sweden", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Instituto de Astrofísica e Ciências do Espaço, Universidade do Porto, CAUP, Rua das Estrelas, 4150-762 Porto, Portugal; Departamento de Física e Astronomia, Faculdade de Ciências, Universidade do Porto, Rua do Campo Alegre, 4169-007 Porto, Portugal", "CNRS/Univ. Bordeaux, LAB, UMR 5804, F-33270, Floirac, France" ]
[ "2016arXiv161104064M", "2017A&A...601A..38J", "2017A&A...605A..98C", "2017AJ....153..258B", "2017AN....338..442A", "2017ApJ...838...85J", "2017ApJ...838...90O", "2017ApJ...839...94B", "2017ApJ...848...34H", "2017ApJ...848...83L", "2017JGRE..122..124N", "2017MNRAS.467.2845A", "2017MNRAS.470.4363C", "2017MNRAS.472.2517J", "2017arXiv171204944H", "2018A&A...618A..78H", "2018ASSP...49..225A", "2018ApJ...853...83H", "2018ApJ...860..109G", "2018ApJ...865...68B", "2018ApJS..237...18K", "2018ApJS..237...38B", "2018IAUS..330..203B", "2018MNRAS.473.5066T", "2018NatAs...2..297U", "2018arXiv180902660M", "2019A&A...629A..80H", "2019AJ....157....7C", "2019AJ....158...61B", "2019ARA&A..57..571J", "2019ApJ...871...63M", "2019ApJ...880...49H", "2019MNRAS.486.2075B", "2019MNRAS.488.5594P", "2020A&A...634A..29J", "2020A&A...639A..66M", "2021A&A...647A..53A", "2021A&A...650A.110M", "2021A&A...650A.194G", "2021A&A...654A.118S", "2021AJ....162..291J", "2021ApJ...920...94E", "2021ApJ...921...95Z", "2021Eleme..17..235P", "2021JGRE..12606731B", "2021MNRAS.501.4596G", "2021MNRAS.502.3704A", "2021MNRAS.504.4968C", "2021MNRAS.504.5788R", "2021Univ....7..118T", "2021plha.book.....K", "2022A&A...658A.194P", "2022A&A...660A..85J", "2022AJ....164...87K", "2022AJ....164..256H", "2022ApJ...927...31T", "2022MNRAS.513.5829W", "2023A&A...676A..52T", "2023ApJ...948...53S", "2023ApJ...956..113D", "2023ApJS..267...18S", "2023BAAA...64...71Z", "2024AJ....167..161K", "2024ApJ...961L..24G", "2024MNRAS.527.3139Z", "2024arXiv240415422H", "2024arXiv240415427G", "2024arXiv240505324H", "2024arXiv240608630M" ]
[ "astronomy" ]
16
[ "stars: abundances", "stars: individual: HD 361", "HD 10700", "HD 121504", "HD 202206", "techniques: spectroscopic", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.3847/0067-0049/226/1/4", "10.48550/arXiv.1607.03130" ]
1607
1607.03130_arXiv.txt
\label{introduction} Like many stars, the Sun is composed overwhelmingly of H and He ($\sim 98.5\%$ by mass), with the elements O, C, Fe, Ne, Si, N, Mg, S and the rest, in descending order, comprising the remainder \citep{Lodders:2009p3091}. How much of a star's mass is in non-H/He elements (termed `metals') and the proportions of the metals are very important parameters to constrain. Elemental abundances help astronomers constrain stellar ages \citep{bond_2010_aa, Nissen15} and trace Galactic chemical evolution \citep{Timmes:1995p3197,Venn:2004p1483,Soubiran:2005p1496}. Most obviously, the abundances of elements in a star---and therefore its protoplanetary disk---help determine the composition of planets that form around that star. In disks with varying chemical make-ups, different minerals will condense out of the gas, creating unique proportions of solids to form planets \citep{Bond:2008p2099, bond_2010_aa}. Systems with more C than O could potentially form planets not of silicates but rather SiC \citep{Kuchner05, bond_2010_aa}. That being said, it is worth noting that detailed observations show little variation from the solar value ${\rm C}/{\rm O} \sim 0.54$ (\citep{Lodders:2009p3091}, such that most stars have ${\rm C}/{\rm O} < 0.8$ \citep{Nissen14, Teske14}. More subtle mineralogical effects could be more common. The Mg/Si ratio in a planet can substantially change the mineral assemblage and mantle viscosity \citep{Umemoto06, Ammann11} and composition \citep{Santos15}. The widespread idea of chemical tagging, among stars with comparable dynamics, seeks to identify a subset of stars with very similar element abundance patterns. In this way, there is potential to find stars that formed in the same stellar cluster \citep[e.g.][]{Mitschang14, Barenfeld13}, in particular stars that may have formed with the Sun \citep[e.g.][]{GonzalezHernandez:2010p7714,Ramirez14b,Nissen15}. Chemical tagging studies benefit from measuring many elements, but especially those elements that vary the most from star to star \citep{Ramirez14b}. Using 11 stars with both physical properties and abundance patterns similar to the Sun, or ``solar twins", \citet{Melendez09} looked for patterns among those refractory elements expected to be trapped in planets. They inferred that the Sun appears to be slightly depleted, relative to these solar twins, in elements trapped within the orbiting planets \citep{Adibekyan12}. Comparatively, \citet{GonzalezHernandez:2010p7714} argued that they did not find such a discrepancy between solar twins that do and do not host planets. For all these reasons it is important to carefully measure elemental abundances in stars \citep{Truitt15}. Of all the non-H/He elements, the abundance of Fe is most easily measured in a star's spectrum, owing to a large number of absorption lines at optical wavelengths. Its abundance (number or molar abundance) in a star's atmosphere is usually reported normalized with respect to the abundance of H in the Sun, expressed as $[{\rm Fe}/{\rm H}] = \log_{10}\left[ ({\rm Fe}/{\rm H}) / ({\rm Fe}/{\rm H})_{\odot}\right]$. The iron-content of nearby, disk stars in the Milky Way spans a range, but the vast majority lie between $-0.5 < [{\rm Fe}/{\rm H}] < +0.5$, meaning stars have between 1/3 and 3 times the Sun's abundance of Fe \citep[e.g.][]{Casagrande11}. When the Galaxy formed, $\sim 1$ Gyr after the Big Bang, the only dominant elements were H and He. But as the Galaxy evolved over time, supernovae, novae, asymptotic giant branch (AGB) stars, etc., enriched the interstellar medium with various elements. The range of metallicities is largely understood as reflecting the time of a star's formation. Indeed, some of the oldest known stars, with metallicity $< -2.0$ dex, illustrate that the ${\rm C}/{\rm O}$ has not been constant over Galactic history \citep{Frebel15}. Given the variety of nucleosynthetic sources occurring over the Galaxy's lifetime, there is no reason to expect that both C and O, or any element, would be produced in exact proportion to Fe. It is widely recognized that in the Galaxy the ratio of $\alpha$-elements (intermediate mass elements with even atomic number Z, as would be produced by successive $\alpha$ captures on elements from carbon up to the iron-peak) compared to Fe shows a trend towards super-solar values at low [Fe/H] \citep[e.g.][]{Timmes:1995p3197}. This is a result of differing contributions from core collapse and Type Ia supernovae over time. What has not been widely understood before the modern era of large surveys is the amount of variation of individual elements in stars of similar [Fe/H]. In other words, while the iron-content is easily measured and is even colloquially synonymous with total metallicity, it only broadly constrains the abundances of other elements. These other elements must be measured and understood in stars. Many research groups have attempted to measure multiple elements, not just Fe, in stellar atmospheres (including many of the authors of this paper). Spectroscopic abundance data for 50 elements across $> 3000$ stars were compiled from 84 literature sources by \citet{Hinkel14}, who created the {\it Hypatia Catalog}. A surprising result of this compilation, similar to that seen in \citet{Torres12}, was that different stellar spectroscopy groups infer quite discrepant abundances for the same elements in the same stars. The variations in elemental abundances often exceed (by factors of 3 or more) the formal uncertainties in the measurements. Discrepancies between varying techniques were large enough, and found often enough, that the lack of agreement represents a crisis in the field of abundance determination by stellar spectroscopy. Indeed, a number of recent papers have sought to compare different abundance methodologies in order to understand their inherent variations. For example, \citet{Bruntt10a} compared a variety of direct and indirect techniques used to calculate the fundamental parameters of bright, solar-like stars. The Gaia-ESO \citep{Gilmore12, Randich13} survey team has also analyzed a combination of many methods to remove and understand systematic differences in stellar abundances. \citet{Smiljanic14} conducted a spectroscopic analysis employing 13 different techniques who performed abundance measurements on 1300 FGK-type stars. In addition, \citet{Jofre14, Heiter15a} and \citet{Jofre15} determined the iron-content, effective temperatures and surface gravities, and the abundances for $\alpha$- and iron-peak elements, respectively, for 34 ``Benchmark Stars'' selected as the cornerstone for the Gaia-ESO survey's data calibration. A variety of tests were performed on the spectra of these stars in order to understand the effect that, for example, data resolution, instrument, and local thermodynamic equilibrium (LTE) approximations may have on different techniques used to analyze the data. In an attempt to reconcile the discrepancies, a Workshop Without Walls called ``Stellar Stoichiometry" was held at Arizona State University, April 11-12, 2013, supported by the NASA Astrobiology Institute. The effects of elemental abundance variations on planet formation, planetary structure, and habitability were discussed \citep{Young14}, but the emphasis was on resolving the abundance inconsistencies. To that end, workshop attendees from stellar spectroscopy groups were asked to participate before and after the workshop in a comparative study (affectionately called our ``homework assignment" but what we will refer to here as the ``Investigation"), in which they were provided a common set of 4 stellar spectra and asked to derive the abundances of several key elements in a variety of ways. This paper reports the results of the comparative study, which was elaborated upon due to the findings and discussion from the initial workshop. In \S \ref{s.homedescription} we describe the comparative study: we list the participants and the numerical methods that each group used and then describe the common data set each group was asked to analyze. In \S \ref{s.results} we present the results from the study, first discussing the stellar parameters and then we present the abundances of 10 key elements inferred by each group. In \S \ref{s.litcomp}, we report on updated findings from the {\it Hypatia Catalog} exemplifying the magnitude of the discrepancies in abundance determinations. We also compare the results of the analysis with respect to literature abundances. In \S \ref{s.disc} we analyze some of the details of the Investigation, namely the error calculations by each group and the effect of varying line lists on the abundance results. We compare results between methods that used the curve-of-growth technique and those that used spectral fitting. We also offer an overall recommendation to the field in order to establish more consistent results between future stellar abundances determinations. Finally, in \S \ref{s.summary}, we summarize the findings of our Investigation.
\label{s.disc} It has been the goal of this research to better understand the abundance measuring techniques between groups and attempt to resolve the inherent issues. However, we wished to make it clear that it has not been our objective to try to determine which group determines the most ``accurate" stellar abundance, or who is closest to calculating the actual amount of an element within a stellar photosphere. While we agree that this is an important issue, we believe that it is necessary to understand what components of a technique give rise to the observed variations. For now our focus has to be on the precision of the techniques, not the accuracy of their results. This is the first step that needs to be taken such that similar issues don't arise again down the road. And if nothing else, once made to agree, we know if the stellar abundances do not reflect a true representation of the element within the star, these results can be scaled such that they are accurate without any loss of precision. \subsection{Error Analysis} \label{s.error} A major facet of the Investigation was to include participants with a relatively wide variety of abundance determination methods. Not only did we want to see how the techniques compared to each other, but also the ways in which each of the methods were able to adapt to the restrictions placed on the abundance measurements. In this way, we were able to get a more in-depth understanding of the more commonly used techniques, including both their strengths and their limitations. The plots in Figs. \ref{elems1}--\ref{elems2} illustrated the overall abundance variation seen between the groups for all stars and analysis standardizations. However, by employing the metric of data similarity (Table \ref{metric}) and the respective error bars for each element, we glossed over the details of the individual uncertainties. In Appendix \ref{a.people}, each group described not only their technique for measuring stellar abundances but also the way in which their errors were calculated. Overall, the four groups who employed the CoG method (ANU, ASU, Carnegie, and Porto) had a similar uncertainty derivation: combining the errors calculated using a standard deviation of the abundances determined by a line-to-line basis and the effect of stellar parameter changes. In comparison, the spectral fitting groups had vastly different techniques for calculating error. For example, Geneva employed a weighted dispersion and Uppsala used a variety of integrated tools to contrast the data with the synthetic spectrum. All of the individually reported errors can be found in Tables \ref{stellarparams1}--\ref{run4abs}. To more clearly evaluate the error determinations for the element abundances, we have compiled the error calculations in Table \ref{errors} in a variety of different manners. For example, the first set (rows 1-6) shows the average of each group's uncertainty determinations for all four stars over all four analyses with respect to the 10 elements (columns). The second set (rows 7-10) reports errors for each run, averaging together the error bars from all of the groups' measurements for all four stars. The last set (rows 11-14) gives errors for each star, taking the mean of all the uncertainties as reported by every group in every run. Finally, the last row shows the average from all 96 calculations, or the four stars across four analyses by six groups. These last values are similar if not the same as the representative error bars, calculated using both medians and means as discussed in \S \ref{s.results}, given in Figs. \ref{elems1}--\ref{elems2}. We see from the first set of error compilations that ASU has some of the overall lowest uncertainty reports for 9 of 10 elements. While this might imply that ASU's abundance determinations are more accurate, it may alternatively be the case that their uncertainties are simply underestimated. Porto's uncertainties tend to be relatively low, but not in all cases. There were few particularly high or low error determinations for ANU, Carnegie, and Uppsala, though Uppsala's tended to be above average. Geneva report the highest errors in seven cases, where six of them were by a wide margin. In the cases of O and Eu II, Geneva's large errors arose because each element only had one line to measure. Therefore, the error did not come from the standard deviation, but directly from the singly fitted line. In these cases, the single-line error is estimated using the covariance matrix by the least square algorithm, and may be overestimated. In addition, the version of iSpec used in this study is not capable of scanning a small range of wavelength-space in order to determine the peak of the absorption line or do small radial velocity corrections to match the line with the synthetic spectra, as is possible in SME. Therefore, it is susceptible to imperfect wavelength calibration, which clearly affects the determination of individual abundances, as was the case for Ni, Eu II, and to a lesser degree Fe. We discuss these issues within the field in general in \S \ref{s.cali}. We also noted in \S \ref{s.parameters} that the Geneva measurements exhibited opposite trends in their stellar parameter determinations between Runs 1 and 3 as compared with the other groups. It is most likely that the mismatch in wavelength calibration, especially when implementing the standardized list, is responsible for the notably different calculations. Also discussed within \S \ref{s.parameters} were the size of the uncertainties for the stellar parameters calculated by the Uppsala group. These errors were not reflected in the abundances, as discussed in Appendix \ref{a.people} and per Table \ref{errors}. However, Tables \ref{stellarparams1}--\ref{stellarparams3} demonstrate that they are roughly twice those determined by the other groups. The uncertainties determined by the Uppsala group are consistent with other line-by-line abundance uncertainties and with the parameter uncertainties that were determined for non-solar type stars. Their stellar parameter uncertainties are essentially dominated by the line-to-line scatter which is caused by errors in $gf$ values as well as modeling shortcomings, missing blends, and continuum uncertainties. By adopting astrophysical $gf$ values, some of these errors would cancel and the precision, but not accuracy, would have improved. This would have resulted in error bars more closely aligned with the other groups. As is currently given, however, the values and their uncertainty are more representative of what is produced for the Gaia-ESO survey, which emphasize accuracy over. In the second set of averages given in Table \ref{errors}, we find that for the majority of elements, the average stellar abundance error is relatively consistent when cycling from Run 1 through Run 4. There are, however, a few exceptions. For example, the errors for Mg, Al, Si, Ni, and Eu II get worse when implementing the standardized line list, although to varying degrees. However, the opposite was true for C and more predominantly O, where the average errors were halved in Runs 3 and 4 as compared to Runs 1 and 2. The results for Eu II, O, and to some degree Mg may be related to the method in which errors are calculated for elements with one absorption line by each group. The tables for Runs 3 and 4 (Tables \ref{run3abs} and \ref{run4abs}) illustrate how Geneva was not able to calculate the A(O) abundance for any star while ANU and Porto had null results for HD~361 and HD~10700 during Run 4. Also ANU, ANU, and Carnegie had sporadic, if any, results for A(Eu II). Similarly, Porto did not report errors, only abundances, for A(Mg) and A(Eu II). It is clear that many of the groups had difficulty with respect to these single-line elements, since omissions did not occur for any other elements (save that ANU was unable to measure carbon in HD~10700 for any run). However, the uncertainties given in Tables \ref{run3abs}, \ref{run4abs}, and \ref{errors} do not always reflect the lower confidence associated with fewer lines, especially when the lines are weak. The case for the 6156.8 \aaa oxygen line, which had a EW of 4.1 m\aaa, is an excellent example. We therefore make note that special precautions should be made, both when measuring and determining the uncertainties, for elements with only a single absorption line. This is an issue that should be addressed within the community and implemented with more transparency and reproducibility within stellar abundance techniques. Finally, for the last section in Table \ref{errors}, with respect to the average errors associated with the individual stars, the most metal-poor star, HD~10700, typically had the largest error bars. This makes sense given the smaller and weaker lines within the spectrum. The other relatively metal-poor star, HD~361, also had slightly larger associated errors. The converse was not always true for HD~202206, which we expected to achieve the lowest uncertainties given that it is metal-rich. However, the average error for A(Fe) was largest for HD~202206. \subsection{Wavelength Calibration} \label{s.cali} As mentioned previously in the paper, we found a significant mismatch in the provided stellar spectra that resulted in a misalignment between spectral features and wavelength numbers below 5050 \aaa. Prior to sending the data to the participating groups, the spectra was tested using the ASU methodology and the results were comparable to other data and literature sources. The ASU team has also looked into the typical wavelength calibration from other data/telescopes and found the provided spectra to be similarly, if not better, calibrated than other (published) sources. In addition, per our analysis of the abundances produced here with respect to the literature sources and their uncertainties (\S \ref{s.litcomp} and Appendix \ref{a.litcomp}), we find that all results are rather standard for the field. However, while the data provided here was ``typical," the wavelength calibration was not perfect. In general, the EW/CoG methods will find the appropriate line, due to either automatic procedures that have a margin of error in which they can search for the correct spectral feature or because many groups tend to measure the EWs manually. While SME has an option to do small radial velocity corrections, that margin is limited. The version of iSpec used in this study does not implement that kind of correction, so it was more affected by the imperfect wavelength calibration than the rest of techniques. Rather than see spectral fitting techniques, which have a number of benefits as compared to the CoG techniques, at a disadvantage, it makes sense that proper wavelength calibration should become more of a priority within the field. Offset spectra could affect the shape of the absorption lines, especially when taking into account convective shifts (see Appendix \ref{a.field}). In fact, the overall impact that wavelength calibration may have on spectroscopic analysis is not clearly understood and could result in some of the systematic scatter that has been noted here and throughout the field. Therefore, although not specifically tested here, we believe that more precise wavelength calibration could benefit all methodologies for determining stellar abundances. \subsection{Spectral Fitting vs. Curve-of-Growth Techniques} Over the course of our analysis, we compared different levels of standardization throughout the analyses, the absolute abundance results between elements, variations seen in stars of assorted spectral types, and calculations from separate abundance techniques. With respect to the stellar parameters, namely Figure \ref{params} and Tables \ref{stellarparams1}--\ref{stellarparams3}, we did not find any correlations between the groups who used spectral fitting (Geneva and Uppsala) and those who used the CoG technique (ANU, ASU, Carnegie, and Porto). In fact, in terms of stellar parameters, the two spectral fitting groups tended to be at the opposite end of the ranges both in terms of values calculated and size of errors reported. No systematic difference was found between groups using the MARCS and ATLAS9 atmosphere models. While EW measurements did affect the results for the CoG techniques, there was no systematic difference between different EW measuring tools. There were a few instances of similarity for the spectral fitting groups when determining the stellar abundances. For example, both Geneva and Uppsala were often at the extreme ends of the spread, either the highest or lowest, when measuring A(Mg) in general. In addition, both groups had the largest determinations for most A(Al) calculations compared with the other groups, while they had lowest results in general for A(Fe) and A(Eu II). However, we note that (1) the individual and representative error bars associated with these elements often overlap with the CoG groups, (2) all measurements for Eu II are inconsistent (or unmeasurable) given the weakness of the line, and (3) there are counterexamples prevalent in each of the above elements. Therefore, we are left to conclude that correlations between the spectral fitting and CoG techniques are more likely coincidental and are not strongly related to the differing methods. \subsection{Varying Line Lists} As discussed in \S \ref{s.elemlinebyline}, during both analyses where the groups could choose which lines they measured, a variety of different line lists were used, especially with respect to iron (see Tables \ref{anulines}-\ref{uppsalalines}). However, iron is one of the most important factors when determining stellar abundances, since it influences stellar parameter calculations and the measurements of other elements. With so few overlapping lines between the groups in this study, it was hard to pinpoint areas of variation. Many line lists used atomic parameters from laboratory or observational sources, such as listed VALD \citep{Kupka00}. However, we found that the atomic parameters, particularly the value of the oscillator strengths, were correlated with abundance discrepancies. The issue with atomic parameters is well known \citep[e.g.][]{Thevenin1999, Doyle13} and often corrected by via a differential analysis technique. This involves using a benchmark star, typically a solar spectrum, to redetermine the $\log(gf)$ values. In this way, an inconsistent atomic parameter can be corrected by reconciling an over- or under-estimated abundance calculation in the solar spectrum, which is then removed from the final result. Differential analysis is a common practice, as seen in \citet[][Table 3]{Hinkel14} and Table \ref{update}, but can be difficult to compare with results that do not normalize to the same solar abundance scale. \citet{Sousa14} tackles the issue of a homogenous line lists specifically with respect to stellar parameter calculations. They found that a consistent line list would improve the accuracy of parameter results, but would require a selection of the ``best" iron lines and agreement within the community. With a variation in the number of lines, the particular lines, and the atomic parameters for Runs 1 and 2, we dealt with many free variables in the calculation of stellar parameters. As shown in the last column Table \ref{models}, the number of Fe I lines varied from 60--298 between the six groups. Yet, in general, the ranges in stellar parameters became worse when implementing the standardized line list, as compared to the autonomous run and shown in Tables \ref{stellarparams1}--\ref{stellarparams3}. However, as discussed in \S \ref{s.elements} and shown in Tables \ref{run1abs}--\ref{metric}, the abundances improved when standardizing the stellar parameters and, to a lesser extent, the line list. Therefore, it would appear that perhaps a line list could be chosen such that the stellar parameters are consistent and precise, which would therefore allow the elemental abundances to be more consistent between groups, such as described in \citet{Smiljanic14}. \citet{Pagano16} test the amount of variation in stellar parameters and final abundance determinations as a function of the number of Fe lines in the linelist using spectra from the same survey used in this work. They find that above $\sim$70 lines parameters and abundances vary by less than the error, with little improvement at higher numbers of lines. For smaller linelists the determinations could change substantially with number of lines and with lines chosen. \subsection{Consistency} In this study, we have attempted to illuminate some of the issues plaguing the stellar abundance community and determine a solution to rectify those differences. We identify quantities that give rise to significant variation and recommend directions for further research. We do not provide specific recommendations of lines, techniques, or parameters to use. Determinations of best practices require a great deal of further investigation. In a very broad sense, we saw that standardizing either the stellar parameters, the line list, or both typically, but not universally, minimized the discrepancies between groups. Whether a fixed line list or fixed stellar parameters were more beneficial changed on an element-by-element basis. By directly comparing the values for measured EWs on a line-by-line basis, it was found that some discrepancies could be attributed to differences in the EW measurements. However, some lines resulted in discrepant abundances even when the EWs were similar or identical. In such cases the atomic constants, in particular oscillator strength, were often found to be the culprits. Therefore, it is our recommendation that a standard group of benchmark stars (for example \citealt{Jofre14, Blanco14a, Heiter15a, Jofre15}) be established to calibrate abundance finding techniques in order to produce consistency in stellar abundances regardless of measurement technique. Any group could adjust their procedure to reproduce the benchmark EWs (if measured), stellar parameters, and abundances. The community may then directly compare the group's abundances using the benchmark technique, which should be consistent across all groups, to the group's preferred method if they believe it to be more accurate. It would also benefit the community to agree upon the best elemental absorption lines and atomic parameters, which would also yield more reliable results. In this way, we may systematically reduce the variation between methods. While we believe our recommendation is reasonable and consistent with our results, we recognize that this will not solve all of the inherent issues. For example, it was our hypothesis at the beginning of this endeavor that stellar parameters and line lists would account for most of the discrepancies, and increasing standardization of these components would improve the agreement of the results. Namely, if identical stellar parameters and line lists were used, the abundances would be very similar and the results in Run 4 would be best. However, we saw throughout the paper that this was not always the case and particular methods or elements were adversely affected. The overall implication is that there are other aspects of the abundance measurement techniques that need to be addressed or standardized, since the remaining scatter is largely unexplained. In addition, elements with only one absorption line, blends, non-LTE effects, HFS, et cetera, require special treatment in order to attain even a hint of uniformity, for example O, Mg, and Eu II. We strongly urge the community to better understand their abundance techniques and the ways in which they may break down, such that these special cases may be properly handled. We also seek to attain error determinations that truly reflect the uncertainties within not only the elements but also the stellar parameters. The error bars may be calculated with respect to benchmark stars, which would ideally provide a basis on which to evaluate abundance variations between groups on a large scale, for example the work conducted by \citet{Smiljanic14}. To date, the results of multiple analyses of the same targets do not produce the same results. Therefore, we call for transparency in presented results to unravel the systematic offsets that are clearly present in our current stellar abundance data. In other words, it would be beneficial if groups published their full line lists, the EW measurements for their lines, atomic parameters, stellar parameters, error bars, and methods of measurement and error calculation in detail. We also believe that it is important to report absolute abundances unnormalized to solar values wherever possible. While this clearly results in longer papers, as demonstrated here and in the Appendices, we believe that it is an important step towards resolving the issues prevalent in the field. {\bf However, transparency of techniques alone is not sufficient if the community does not come together to increase reproducibility. } In this vein, we believe that more rigorous comparisons between data sets should be implemented, specifically those which involve statistics, and an avoidance of purely graphical comparisons. We also encourage members of the field to openly discuss offsets and variations found between techniques in order to discover their fundamental root. Since there is a current lack of ``ground-truthing" within astronomy, methodologies are not so much ``wrong" as incomparable, which has lead to the current lack of consensus. Once abundance results are accurate and agree between techniques, only then can we focus on determining the true, precise stellar chemical make-up. We are optimistic that stellar abundances can become both precise and accurate, however, it is clear that it requires an effort that spans the entire community. As a result of the preliminary discussion and tests at ASU's Stellar Stoichiometry workshop, we have formulated an investigation to understand the causes of variations between stellar abundance measurement techniques. A total of six groups were asked to determine the abundances within four stars, exhibiting a range in iron-content, over a series of multiple tests. Each group was asked to calculate the abundances for each star a total of four times, where Run 1 utilized their own autonomous method, Run 2 employed standardized stellar parameters, Run 3 implemented a standardized line list, and Run 4 had both standardized stellar parameters and element absorption lines. The results for the stellar parameters and the 10 elements, namely C, O, Na, Mg, Al, Si, Fe, Ni, Ba II, and Eu II, can be found in Tables \ref{stellarparams1}--\ref{run4abs}. The effect of the standardized line list on the stellar parameters appeared to have a polarizing effect, as summarized in Table \ref{sum}. In some cases, using the same lines allowed measurements to become more consistent between groups. On the other hand, a number of the stellar parameter values became more disparate during Run 3. These variations were not due to whether the methodology employed spectral fitting or CoG. However, the cause may be related to the inherent degeneracy between the stellar parameters or the number of Fe lines employed by each of the groups. It was our expectation at the beginning of the Investigation that Run 4 would yield the most comparable results between groups. While the spread in the abundance determinations did decrease during this run as compared to Run 1, the results were not as consistent as we had expected. We implemented a metric of data similarity (Table \ref{metric}) in order to quantify the similarity/dispersion of the abundance determinations. In general, we found that Run 2 had consistently better results between the elements and stars, followed by Run 4. As seen in Table \ref{sum}, out of the 40 possibilities (10 elements within 4 stars), using the standardized parameters (only) achieved the most similar measurements between groups a total of 16 times. Using both the standardized parameters and lines resulted in the most comparable abundances in 12 instances. Given that both Runs 1 and 3 had similar success rates for best abundance similarity for a star between analyses, the implication is that the standardized line list may have adversely affected the abundance techniques for particular elements. This conclusion is mirrored in our line-by-line abundance analysis with respect to iron. Unfortunately, the small number of non-Fe lines measured by all groups for all four analyses made comparisons and broader conclusions difficult. Given that Run 3 generated much larger spreads in stellar parameters than Run 1, it may be that a standardized line list could improve consistency, but the choice of Fe lines used must be optimized. On the whole, elements that had a strong dependence on stellar parameters, as measured by an improvement between Run 1 to Run 2, did not improve or get worse during Run 3. We saw similar spreads in the data when comparing the values determined here to those found within the literature. Optimistically, this shows that the setup of the Investigation accurately reflected the techniques and measurements found within the community. In addition, it emphasized the point that particular elements, such as C and O, require special attention to determine precise abundances. For elements like Fe and Al, and to a lesser extent Na, Al, Si, and Ni, the overall abundance determinations may be sensitive to either the lines, the technique, or the stellar parameters used in their calculation. The Investigation that was conducted here is complimentary to studies by \citet{Smiljanic14, Jofre14, Jofre15}, all of which are relevant to ongoing and upcoming large stellar surveys, such as the Gaia-ESO survey, the Transiting Exoplanet Survey Satellite (TESS), the Characterizing Exoplanets Satellite (CHEOPS), and the James Webb Space Telescope (JWST). It is important that stars be precisely and consistently measured such that they, and possible orbiting exoplanets, may be well characterized. It was our hope that either standardized stellar parameters, line list, or both would ameliorate the disparate abundance measurements between different techniques. Instead, we have found a deeper understanding of the problems inherent to the stellar abundance field. The Investigation has shown that, in order for measurements to be copacetic, 1) the details of the methods and their input parameters need to be directly compared, 2) particular elements require special care, 3) line lists and/or stellar parameters should be standardized, and importantly, 4) all research should be presented with the utmost transparency to ensure reproducibility of the results. Given the huge amount of determination, flexibility, support, and feedback offered by all of the participating groups, we are hopeful that the community can work together to rectify these issues and work towards not only precise but accurate stellar abundances.
16
7
1607.03130
Stellar elemental abundances are important for understanding the fundamental properties of a star or stellar group, such as age and evolutionary history, as well as the composition of an orbiting planet. However, as abundance measurement techniques have progressed, there has been little standardization between individual methods and their comparisons. As a result, different stellar abundance procedures determine measurements that vary beyond the quoted error for the same elements within the same stars. The purpose of this paper is to better understand the systematic variations between methods and offer recommendations for producing more accurate results in the future. We invited a number of participants from around the world (Australia, Portugal, Sweden, Switzerland, and the United States) to calculate 10 element abundances (C, O, Na, Mg, Al, Si, Fe, Ni, Ba, and Eu) using the same stellar spectra for four stars (HD 361, HD 10700, HD 121504, and HD 202206). Each group produced measurements for each star using (1) their own autonomous techniques, (2) standardized stellar parameters, (3) a standardized line list, and (4) both standardized parameters and a line list. We present the resulting stellar parameters, absolute abundances, and a metric of data similarity that quantifies the homogeneity of the data. We conclude that standardization of some kind, particularly stellar parameters, improves the consistency between methods. However, because results did not converge as more free parameters were standardized, it is clear there are inherent issues within the techniques that need to be reconciled. Therefore, we encourage more conversation and transparency within the community such that stellar abundance determinations can be reproducible as well as accurate and precise.
false
[ "HD", "data similarity", "Stellar elemental abundances", "different stellar abundance procedures", "abundance measurement techniques", "individual methods", "absolute abundances", "methods", "such that stellar abundance determinations", "the same stellar spectra", "the resulting stellar parameters", "inherent issues", "a standardized line list", "little standardization", "the same stars", "evolutionary history", "results", "both standardized parameters", "recommendations", "Eu" ]
8.50328
10.13675
-1
12510344
[ "Gusakov, M. E.", "Dommes, V. A." ]
2016PhRvD..94h3006G
[ "Relativistic dynamics of superfluid-superconducting mixtures in the presence of topological defects and an electromagnetic field with application to neutron stars" ]
30
[ "Ioffe Physical-Technical Institute of the Russian Academy of Sciences, Polytekhnicheskaya 26, 194021 Saint-Petersburg, Russia; Peter the Great Saint-Petersburg Polytechnic University, Polytekhnicheskaya 29, 195251 Saint-Petersburg, Russia", "Ioffe Physical-Technical Institute of the Russian Academy of Sciences, Polytekhnicheskaya 26, 194021 Saint-Petersburg, Russia" ]
[ "2017CQGra..34l5002A", "2017JApA...38...43C", "2017MNRAS.467L.115D", "2017MNRAS.469.3928K", "2017MNRAS.469.4979P", "2017MNRAS.471..507C", "2017PhRvD..96j3012G", "2018ASSL..457..455S", "2018PhRvD..98d3007O", "2019JPhCS1400b2006G", "2019MNRAS.482.2573D", "2019MNRAS.485.4936G", "2020CQGra..37b5014G", "2020MNRAS.493..382S", "2020MNRAS.498.3000C", "2020MNRAS.499.4561G", "2020NuPhB.95515049D", "2020PhRvD.101j3020D", "2020PhRvD.102f3011R", "2021MNRAS.503.1407S", "2021MNRAS.506L..74K", "2021PhRvD.104l3008D", "2021Univ....7...17A", "2021Univ....7...28G", "2022EPJA...58...42S", "2022atcc.book..219A", "2023MNRAS.520.1872B", "2023MNRAS.521.5724T", "2023PhRvC.108b5814G", "2024MNRAS.527.9431M" ]
[ "astronomy", "physics", "general" ]
3
[ "General Relativity and Quantum Cosmology", "Astrophysics - High Energy Astrophysical Phenomena", "Condensed Matter - Superconductivity" ]
[ "1956RSPSA.238..215H", "1960AdPhy...9...89H", "1967rhm..book.....L", "1969Natur.224..673B", "1969Natur.224..674B", "1971ApJ...168..175W", "1975UsFiN.116..413G", "1977PhRvD..16..275E", "1981JETP...54..919V", "1982JETP...56..923L", "1982PhLA...91...70K", "1984ApJ...282..533A", "1987PhRvA..36.3947H", "1991AnPhy.205..110M", "1991ApJ...380..515M", "1991ApJ...380..530M", "1995ApJ...447..305S", "1995NuPhB.454..402C", "1996PhRvC..54.2745B", "1997MNRAS.290..203S", "1998MNRAS.296..903M", "1998MNRAS.297.1189L", "1998NuPhB.531..478C", "1998clel.book.....J", "2000PhRvB..62.9740C", "2001MNRAS.325..426K", "2001PhRvD..65a4022I", "2002PhRvD..66a4015I", "2004ARA&A..42..169Y", "2004PhRvD..69d3001P", "2005NuPhA.761..333G", "2005PhRvB..72u2508P", "2005PhRvD..71h3003S", "2005PhRvD..71h3006P", "2005qvhi.book.....D", "2006MNRAS.372.1776G", "2006PhRvC..73d5802C", "2007PhRvD..76h3001G", "2008MNRAS.383.1551A", "2008PhRvD..78h3006G", "2009PhRvC..79e5806G", "2009PhRvC..80a5803G", "2009PhRvD..79d3004K", "2010JPhG...37g5202A", "2010PhRvC..81b5804G", "2011MNRAS.410..805G", "2011PhRvD..83j3008K", "2012PhRvD..86f3002H", "2013MNRAS.428.1518G", "2013MNRAS.428L..26G", "2013MNRAS.434..123V", "2013arXiv1302.6626P", "2014MNRAS.439..318G", "2014MNRAS.442.3484K", "2014arXiv1406.6109G", "2015IJMPD..2430008H", "2015MNRAS.453..671G", "2015PhRvC..91c5805S", "2016CQGra..33x5010A", "2016MNRAS.455.2852D", "2016PNAS..113.3944G", "2016PhRvD..93f4033G" ]
[ "10.1103/PhysRevD.94.083006", "10.48550/arXiv.1607.01629" ]
1607
1607.01629_arXiv.txt
Assume that we have a relativistic magnetized finite-temperature plasma (possibly in the strong gravitational field) composed of superfluid neutral particles, superconducting positively charged particles and normal (nonsuperconducting) negatively charged particles. Depending on the density, the positively charged particles may form either type-I or type-II superconductor, and the plasma may contain topological defects -- Feynman-Onsager and/or Abrikosov vortices. What are the macroscopic dynamic equations describing such a system? The question is not so far-fetched as it may seem at first glance. For example, the neutron-proton-electron ($npe$) mixture in the outer neutron-star cores meets all the conditions formulated above. First, it is relativistic and magnetized. The typical surface magnetic field is $B \sim 10^8 \div 10^{15}$~G \cite{vigano_etal_13, hpy07} and is likely to be larger in the deeper layers \cite{gwh16}; the surface gravitation acceleration is also huge, $g_{\rm s} \sim 2 \times 10^{14}$~cm~s$^{-2}$ \cite{hpy07}, electrons are ultra-relativistic, while neutrons can be moderately relativistic. Second, according to microscopic calculations \cite{gps14, ls01}, confirmed (to some extent) by observations of cooling and glitching neutron stars \cite{yp04, plps13, hm15}, neutrons and protons in their interiors become superfluid/superconducting at temperatures $T\lesssim T_{{\rm c}i}$, where $T_{{\rm c}i} \sim 10^8 \div 10^{10}$~K is the nucleon critical temperature ($i=n$, $p$). Third, in a rotating magnetized neutron star it can be energetically favourable to form Feynman-Onsager/Abrikosov vortices \cite{sauls89} (the latter are formed only if protons are type-II superconductor; if, instead, they are of type-I, different structures appear, see Sec.\ \ref{TypeI} for more details). Thus it is not surprising that the dynamic properties of magnetized superfluid-superconducting neutron-star plasma have been the subject of numerous studies in the past, both in nuclear matter (see, e.g., Refs.\ \cite{ep77, als84, vs81, hk87, ml91, mendell91a, mendell91b, ss95, mendell98, gas11, prix05}) and in quark matter (e.g., Refs.\ \cite{ib02a, ib02b, as10, hac12}). In particular, Vardanyan and Sedrakyan \cite{vs81} were the first who generalized hydrodynamics of a mixture of two superfluids \cite{khalatnikov00,ab76} to charged superfluids coupled to the electromagnetic field. These equations were further extended by Holm and Kupershmidt \cite{hk87} to $N$ charged superfluids, who derived these equations from the Hamiltonian formalism. Finally, the most general {\it nonrelativistic} finite-temperature equations, describing charged superfluids and accounting for the mutual friction forces \cite{hall60,hv56} between various liquid components, were formulated by Mendell and Lindblom \cite{ml91}, who used in their work the ideas of Refs.\ \cite{bk61, khalatnikov00, hk87}. This important work was subsequently used by Mendell \cite{mendell91a,mendell91b} who applied the equations of Ref.\ \cite{ml91} to neutron stars, assuming that all neutrons and protons are paired (i.e., $T\ll T_{{\rm c}i}$). (A little bit later, Sedrakian and Sedrakian \cite{ss95} did a similar job by extending the results of Ref.\ \cite{vs81} to include dissipation and mutual friction forces in their equations.) In his work, Mendell formulated a set of simplified magnetohydrodynamic equations, but, unfortunately, incorrectly identified the magnetic field ${\pmb H}$ with the magnetic induction ${\pmb B}$ and the electric displacement ${\pmb D}$ with the electric field ${\pmb E}$. The first of these inaccuracies (identification of ${\pmb H}$ with ${\pmb B}$) was noticed in Ref.\ \cite{cpl00} and corrected by Glampedakis, Andersson, and Samuelsson \cite{gas11} (hereafter GAS11); the second inaccuracy (identification of ${\pmb D}$ with ${\pmb E}$) is discussed here (see Appendix \ref{vortapp}). Except for the corrected inaccuracy, the GAS11 version of magnetohydrodynamics is equivalent (up to notations) to that of Mendell \cite{mendell91a} and is the most advanced treatment of superfluid-superconducting mixtures in neutron stars up to date. It is derived using the variational framework \cite{prix04,prix05} and assuming $T=0$. All the works discussed by us so far were performed in the nonrelativistic approximation. This is a rather serious shortcoming because, as we have already mentioned, neutron stars are essentially relativistic objects. The extension of magnetohydrodynamics of GAS11 (as well as more general equations of Ref.\ \cite{ml91}) to the relativistic case is not trivial. For uncharged one-component superfluids this problem has been addressed in Refs.\ \cite{kl82,lk82,lsc98,carter00,cl95,kg12,dg16,awv16} and has recently been ``solved'' in Ref.\ \cite{gusakov16} (hereafter G16). We are aware of only one attempt \cite{cl98} to consider charged mixtures in full relativity. This reference neglected all dissipation effects (including mutual friction) and studied only the low-temperature case $T\ll T_{{\rm c}i}$; unfortunately, it did not provide a nonrelativistic limit for the derived equations so that it is hard to compare them with the formulations available in the literature. Note that Ref.\ \cite{cl98} adopted the variational approach similar to that developed in Ref.\ \cite{cl95} in application to uncharged superfluids. This approach was criticised in G16 (see Appendix F there) where it was argued that it does not reproduce the well established nonrelativistic Hall-Vinen-Bekarevich-Khalatnikov superfluid hydrodynamics \cite{bk61, khalatnikov00}. We believe the same conclusion applies also to the results of Ref.\ \cite{cl98}. The aim of the present study is to fill the existing gaps and derive a set of relativistic finite-temperature equations describing superfluid-superconducting mixtures, bearing in mind application of these results to magnetized rotating neutron stars. As in Refs.\ \cite{bk61} and G16, our derivation rests on the consistency between the conservation laws and the entropy generation equation. For definiteness, in this paper we consider a liquid composed of superfluid neutrons ($n$), superconducting protons ($p$), and normal electrons ($e$). Extension of our results to more complicated compositions is straightforward (see, e.g., Refs.\ \cite{gk08, kg09, dg16, hac12}). Here we are mostly interested in the non-dissipative equations (but we allow for mutual friction dissipation, see Remark~1 in Sec.\ \ref{TypeII}). Correspondingly, we assume that neutron and proton thermal excitations as well as electrons move with one and the same ``normal'' four-velocity $u^{\mu}$. In what follows all thermodynamic quantities are defined in the frame comoving with the normal (nonsuperfluid) liquid component, in which $u^{\mu}=(1,0,0,0)$. By default, any 3d-vector appearing in the text (e.g., magnetic induction ${\pmb B}$) is written in that frame. The paper is organized as follows. Section \ref{Maxwell} introduces Maxwell's equations in the medium written both in the standard and explicitly Lorentz-covariant form. Section \ref{no} considers uncharged and charged mixtures in the absence of vortices and other magnetic domain structures. In Sec.\ \ref{setup} we discuss the strategy for generalization of equations of Sec.\ \ref{no} in order to allow for the topological defects and related bound charges and currents in the mixture. In Sec.\ \ref{TypeI} this strategy is applied to derive the corresponding dynamic equations under assumption of type-I superconductivity of protons. Section \ref{TypeII} is devoted to considering type-II proton superconductivity and accounting for the possible presence of both neutron (Feynman-Onsager) and proton (Abrikosov) vortices. Section \ref{symmetry} proves that the energy-momentum tensors obtained in Secs.\ \ref{TypeI} and \ref{TypeII} are symmetric, and expresses them through a set of phenomenological coefficients which can be calculated by specifying a microscopic model for the energy-density of the mixture. The general dynamic equations of Sec.\ \ref{TypeII} are simplified for typical neutron-star conditions in Sec.\ \ref{MHDapprox}. Finally, we sum up in Sec.\ \ref{summary}. The paper also contains a number of appendices, where we present technical, more model-dependent, or less important results. In particular, Appendix \ref{notation} introduces some basic notation used throughout the paper. Appendix \ref{corresp} provides a correspondence table between our notation and that adopted in G16. Appendix \ref{transformation} contains an example of the energy density transformation used in Secs.\ \ref{TypeI} and \ref{TypeII}. Appendix \ref{Abraham} reveals the relation between the energy-momentum tensor of Sec.\ \ref{TypeI} and the well known Abraham tensor. Appendix \ref{vortex1} discusses some general relations characterizing isolated neutron or proton vortices. Appendix \ref{why} demonstrates that there exist some bound charges associated with each moving vortex. Appendix \ref{deterapp} presents two simple microscopic models allowing one to determine the phenomenological coefficients from Sec.\ \ref{symmetry}. Finally, Appendix \ref{sumapp} contains the full set of dynamic equations derived in Secs.\ \ref{TypeI} and \ref{TypeII}, and Appendix \ref{nonrel} analyses the nonrelativistic limit of simplified equations of Sec.\ \ref{MHDapprox}. Unless otherwise stated, in all sections except for Sec.\ \ref{Maxwell} and Appendices \ref{vortex1}, \ref{why}, \ref{deterapp}, and \ref{nonrel} the speed of light $c$, the Planck constant $\hbar$, and the Boltzmann constant $k_{\rm B}$ are set to unity, $c=\hbar=k_{\rm B}=1$. Throughout the paper we assume that the spacetime metric is flat, $g_{\mu\nu}={\rm diag}(-1,\, 1,\, 1,\, 1)$. Generalization of our results to arbitrary $g_{\mu\nu}$ is straightforward and can be achieved by replacing ordinary derivatives in all equations with their covariant counterparts.
\label{summary} This paper is devoted to studying the dynamic properties of superfluid-superconducting mixtures in neutron stars accounting for the possible presence of electric and magnetic fields, as well as neutron (Feynman-Onsager) and proton (Abrikosov) vortices. Our results and main conclusions are summarized as follows: 1. Using the method and ideas from Refs.\ \cite{bk61} and G16, we derived a set of fully {\it relativistic} equations (see Appendix \ref{sumapp}) describing a charged mixture composed of superfluid neutrons, superconducting protons, and electrons (the simplest neutron-star composition). Generalization of these equations to more exotic compositions (including, e.g., muons, hyperons, etc.) is straightforward \cite{dg16,kg09,gk08, hac12}. 2. The proposed equations can be used at {\it finite} temperatures, i.e., they allow for the possible presence of neutron and proton (Bogoliubov) thermal excitations. This is especially important for a sufficiently hot neutron stars, such as magnetars, whose internal temperatures can be $\sim 10^8$~K, i.e., of the order of the nucleon critical temperatures $T_{{\rm c}i}$ \cite{kkpy14, vigano_etal_13} (we remind that at $T>T_{{\rm c}i}$ nucleon species $i=n$, $p$ is completely nonsuperfluid). 3. The derived dynamic equations are ``nondissipative'' in a sense that to obtain them we assume that normal (nonsuperfluid) liquid components (electrons, nucleon thermal excitations, and entropy) move with one and the same velocity (i.e., diffusion effects are ignored). However, we do take into account the {\it mutual friction dissipation } [see Eqs.\ (\ref{sfleq5}) and (\ref{f3})]. Extension of our results to a fully dissipative problem is rather easy and will be reported elsewhere. 4. Estimates show that protons form type-II superconductor in the outer neutron-star core, but become of type-I in the inner core (e.g., GAS11, \cite{aw08,ssz97,ss15,sedrakyan05}). The dynamic equations are derived and analysed in both these cases with the special emphasis on the more elaborated type-II case. It seems that the dynamics of type-I superconductor is discussed for the first time (in the astrophysical context), but the analysis presented is rather brief and simplified, and should be considered as a first step towards the solution of this complex problem. 5. Our main results include the ``electromagnetic'' energy-momentum tensors $\mathcal{T}_{({\rm E}) }^{\mu\nu}$ (\ref{Te1app}) and $\mathcal{T}_{({\rm M}) }^{\mu\nu}$ (\ref{Tm1app}), and the nucleon ``vortex'' energy-momentum tensors $\mathcal{T}_{({\rm VE}) }^{\mu\nu}$ (\ref{vortexEapp}) and $\mathcal{T}_{({\rm VM}) }^{\mu\nu}$ (\ref{vortexapp}), as well as the ``superfluid'' equations for the cases of type-I (\ref{Vappn}), (\ref{Vappp}) and type-II (\ref{sflrot1app}) proton superconductivities. Remarkably, the vortex energy-momentum tensors have the same structure and are obtained exactly in the same way as the electromagnetic tensors (\ref{Te1app}) and (\ref{Tm1app}) (see Remark 4 in Sec.\ \ref{TypeII}). 6. As a by-product of our work it is shown that for normal matter the sum $\mathcal{T}_{({\rm E}) }^{\mu\nu}+\mathcal{T}_{({\rm M}) }^{\mu\nu}$ of the electromagnetic energy-momentum tensors is directly related to the so called Abraham tensor $T^{\mu\nu}_{\rm Abraham}$ of the standard electrodynamics of continuous media \cite{ginzburg73, gu76, toptygin15}. Thus, our results can be considered as one more derivation of this tensor based on the conservation laws and the requirement that the entropy of a nondissipative closed system remains constant. 7. The equations derived in this paper [in particular, the expressions for the electromagnetic and vortex energy-momentum tensors $\mathcal{T}_{({\rm E}) }^{\mu\nu}$, $\mathcal{T}_{({\rm M}) }^{\mu\nu}$, $\mathcal{T}_{({\rm VE}) }^{\mu\nu}$, and $\mathcal{T}_{({\rm VM}) }^{\mu\nu}$] depend on the four-vectors $E^{\mu}$, $B^{\mu}$, $\mathcal{V}^{\mu}_{({\rm E}i)}$, $\mathcal{V}^{\mu}_{({\rm M}i)}$ and the complementary four-vectors $D^{\mu}$, $H^{\mu}$, $\mathcal{W}^{\mu}_{({\rm E}i)}$, $\mathcal{W}^{\mu}_{({\rm M}i)}$. The physical meaning of these four-vectors is described in detail in the text. For example, the spatial components of $E^{\mu}$, $B^{\mu}$, $D^{\mu}$, and $H^{\mu}$ reduce, respectively, to the electric field, magnetic induction, electric displacement, and magnetic field in the comoving frame moving with the normal liquid component (see Sec.\ \ref{EBDH}); the other four-vectors are related to vortices. The four-vectors mentioned above are not all independent. To express the quantities $D^{\mu}$, $H^{\mu}$, $\mathcal{W}^{\mu}_{({\rm E}i)}$, $\mathcal{W}^{\mu}_{({\rm M}i)}$ through $E^{\mu}$, $B^{\mu}$, $\mathcal{V}^{\mu}_{({\rm E}i)}$, $\mathcal{V}^{\mu}_{({\rm M}i)}$ one should specify, as in the usual electrodynamics of continuous media, the microphysics model for the mixture. This is done, for two simple models, in Appendix \ref{deterapp} (in particular, one of these models analyses the system of noninteracting vortices). However, it is important to point out that the general equations obtained here will likely remain unchanged if one considers more complex models. The only thing that should be modified in the latter case is the relations between the fields $D^{\mu}$, $H^{\mu}$, $\mathcal{W}^{\mu}_{({\rm E}i)}$, $\mathcal{W}^{\mu}_{({\rm M}i)}$ and $E^{\mu}$, $B^{\mu}$, $\mathcal{V}^{\mu}_{({\rm E}i)}$, $\mathcal{V}^{\mu}_{({\rm M}i)}$. 8. It is instructive to compare our results with the most advanced nonrelativistic magnetohydrodynamics of GAS11, describing superfluid-superconducting mixtures. In comparison to GAS11 we: (i) take into account the relativistic and finite-temperature effects; (ii) provide a general framework allowing one to easily incorporate new physics into the existing dynamic equations; and (iii) demonstrate that the electric displacement field ${\pmb D}$ is {\it not} generally equal to the electric field ${\pmb E}$, contrary to what was assumed in GAS11 and some other papers starting from the work by Mendell \cite{mendell91a} (see also Ref.\ \cite{ss95}). 9. The rather complex general system of equations derived in this work can be substantially simplified for typical neutron-star conditions, for which a kind of ``magnetohydrodynamic'' approximation is justified. This approximation is analogous to the usual magnetohydrodynamic approximation for ordinary stars. The corresponding equations are derived and analysed for a simple model of Appendix \ref{vortapp} in Sec.\ \ref{MHDapprox}; their nonrelativistic limit is presented in Appendix \ref{nonrel}, where we also derive a ``magnetic field evolution equation'' (\ref{MagnEvol2}). It is shown that the latter equation coincides with that proposed in Ref.\ \cite{kg01}, but differs from the evolution equation derived in Ref.\ \cite{gagl15} using magnetohydrodynamics of GAS11.
16
7
1607.01629
The relativistic dynamic equations are derived for a superfluid-superconducting mixture coupled to an electromagnetic field. For definiteness, and bearing in mind possible applications of our results to neutron stars, it is assumed that the mixture is composed of superfluid neutrons, superconducting protons, and normal electrons. We analyze the proton superconductivity of both types I and II and allow for the possible presence of neutron and proton vortices (or magnetic domains in the case of type-I proton superconductivity). The derived equations neglect all dissipative effects except for the mutual friction dissipation and are valid for arbitrary temperatures (i.e., they do not imply that all nucleons are paired), which is especially important for magnetar conditions. It is demonstrated that these general equations can be substantially simplified for typical neutron stars, for which a kind of magnetohydrodynamic approximation is justified. Our results are compared to the nonrelativistic formulations existing in the literature, and a number of discrepancies are found. In particular, it is shown that, generally, the electric displacement D does not coincide with the electric field E , contrary to what is stated in previous works. The relativistic framework developed here is easily extendable to account for more sophisticated microphysics models, and it provides the necessary basis for realistic modeling of neutron stars.
false
[ "neutron stars", "typical neutron stars", "superconducting protons", "previous works", "superfluid neutrons", "magnetar conditions", "magnetic domains", "normal electrons", "magnetohydrodynamic approximation", "possible applications", "realistic modeling", "arbitrary temperatures", "neutron and proton vortices", "II", "I", "the proton superconductivity", "the electric field E", "discrepancies", "an electromagnetic field", "the possible presence" ]
5.255332
2.68799
-1
1384972
[ "Peek, J. E. G.", "Bordoloi, Rongmon", "Sana, Hugues", "Roman-Duval, Julia", "Tumlinson, Jason", "Zheng, Yong" ]
2016ApJ...828L..20P
[ "The First Distance Constraint on the Renegade High-velocity Cloud Complex WD" ]
7
[ "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "MIT-Kavli Center for Astrophysics and Space Research, 77 Massachusetts Avenue, Cambridge, MA 02139, USA", "Institute of Astronomy, KU Leuven, Celestijnenlaan 200 D, B-3001 Leuven, Belgium", "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 Astronomy, Columbia University, New York, NY 10027, USA" ]
[ "2017ApJ...840...65Z", "2018MNRAS.474..289W", "2018RNAAS...2...59P", "2019ApJ...881....4S", "2019ApJ...884...53F", "2022MNRAS.509.5756G", "2023MNRAS.522.1535R" ]
[ "astronomy" ]
9
[ "Galaxy: evolution", "Galaxy: formation", "Galaxy: halo", "ISM: atoms", "ISM: clouds", "ISM: kinematics and dynamics", "Astrophysics - Astrophysics of Galaxies" ]
[ "1980ApJ...236..577B", "1991A&A...250..509W", "1995A&A...302..364S", "1997ARA&A..35..217W", "1998ApJ...500..525S", "2000ApJ...544L.107W", "2001ApJS..136..463W", "2003ApJ...597..948P", "2003SPIE.4841.1682P", "2004AJ....127..899S", "2007ApJ...656..907P", "2007ApJ...670L.113W", "2008ApJ...672..298W", "2008ApJ...679L..21L", "2008ApJ...684..364T", "2008ApJ...684.1143X", "2008Msngr.133...17M", "2010A&A...521A..17K", "2010ApJ...722..367F", "2011ApJ...737..103S", "2012ARA&A..50..491P", "2012ApJ...749..181F", "2012ApJ...750...99T", "2012ApJ...756..158S", "2012MNRAS.424.2896L", "2013ApJS..205...20M", "2016ApJ...816L..11F" ]
[ "10.3847/2041-8205/828/2/L20", "10.48550/arXiv.1607.06465" ]
1607
1607.06465_arXiv.txt
High-velocity clouds (HVCs) provide a unique window into the coolest component of the circumgalactic medium and the processes of Galactic inflow and outflow. HVCs, and the complexes into which they are arranged, are found by their emission in \hi or absorption in numerous metal lines, and have radial velocities inconsistent with Galactic rotation \citep{Wakker_1997}. The precise origin of most HVCs is unknown, and some mix of Galactic fountain \citep[e.g.][]{Bregman_1980}, multiphase accretion \citep[e.g.][]{Fern_ndez_2012}, and gas stripping from satellites is typically invoked \citep{Putman_2012}. The exception is the Magellanic stream, which was stripped from the large and small Magellenic clouds, and which we will exclude from our discussion in this work. HVCs with negative radial velocities, which are metal enriched in the range of 10\% to 30\% of the solar metallicity, are likely a tracer of the process by which material accretes onto the Galaxy, though the total rate of this accretion is very uncertain. Less explored are the HVCs with positive radial velocities, most of which are in the inner two quadrants of the Galactic sky. These include the Wannier complexes WA, WB, WD, WE, and the Smith Cloud \citep{1991A&A...250..509W}. The Smith cloud has received significant attention of late, for its strongly cometary appearance which provides enough information to infer past trajectories, and make some inference as to its origin \citep{Lockman_2008, Fox_2015}. Complex WD is the largest-area positive velocity HVC Complex, covering 310 square degrees with a total \hi flux of 1.2 $\times 10^7$ K km s$^{-1}$ arcmin$^2$, and a maximum \hi column density of $\sim 1.2 \times 10^{20}$ cm$^{-2}$. It is by far the largest complex that exists in the inner two Galactic quadrants, where a small fraction of HVC flux is detected. With a range of velocities between +90 and +130 km s$^{-1}$, it is consistent with cylindrical rotation on the far side of the inner Galaxy, 20 kpc from the sun with a mass of 6 $\times 10^7 M_\odot$. This would make it very similar in mass, Galactocentric radius, and height to Complex C, the largest area and brightest HVC complex \citep{Thom_2008}. % One major issue in gaining a better physical understanding of these enigmatic clouds is their unknown distance. Since there are no objects of fixed luminosity in HVCs, there are effectively no intrinsic distance measures. \hi emission or optical and UV absorption lines toward extragalactic background sources only provide distance-independent column densities. HVC distances not only give us a masses for these structures, but also a context; the spatial relationship between the cloud and the nearby spatial and kinematic structure of the disk gives us insight as to its origin. There are a number of indirect methods for measuring the distance to an HVC complex, including H$\alpha$ emission and kinematic structure \citep{Putman_2003, Peek_2007}, but the only proven direct distance measure is stellar absorption. By observing stars with measured distances at medium or high spectral resolution, one can look for absorption lines in Na {\sc i}, \caii H \& K, Ti {\sc ii}, and numerous ultraviolet absorption lines at the velocity of \hi emission from HVCs \cite{1995A&A...302..364S}. By finding detections and non-detections of these absorption lines along lines of sight toward \hi emitting HVCs, distances can be robustly measured. A number of clouds have well-measured distances using this method, but complex WD is not among them \citep{Wakker_2001, Wakker_2007, Thom_2008, Wakker_2008}. In this work we report the first distance upper limit on Complex WD using medium resolution absorption line spectroscopy toward a blue horizontal branch star. We extend the methods of \citet{Sirko_2004} to find the spectral type of the star, and thus put a precise distance limit. We use this to make some inferences as to the possible origin of Complex WD, and how it fits into the structure of Galactic HVCs as a whole.
In this work we have found that while the H$\epsilon$ line is not as powerful a discriminant between BHB, BS, and MS stars as the H$\gamma$ or H$\delta$ line, it can be used to show that USNO-A0600-15865535 is a BHB star approximately 4.4 kpc from the sun. We demonstrated that this star has clear Calcium H \& K absorption lines at a velocity coincident with Complex WD, and that therefore Complex WD must be closer than 4.4 kpc. We used this fact to rule out the originally assumed model of Complex WD, that is a Complex C-like cloud on the far side of the Galaxy, corotating with disk. Furthermore we investigated an intermediate distance scenario in which the Complex resides at $\sim 4$ kpc and found it difficult to reconcile the fixed LSR velocity of the cloud with the strong gradients implied from both the reflex solar motion and differential Galactic rotation. A "near" scenario, wherein a the complex was ejected from the solar vicinity to a distance of a few kpc seemed the most likely, though it would make the HVC the lowest in altitude known. Future observations toward USNO-A0600-15865535 would enable precise determinations of metallicities using other elements, which would help determine whether a disk-origin for this cloud is likely. Further observations towards closer stars, and across the face of the cloud could give us much more detailed information about the distance and three-dimensional morphology of the cloud, which would also help us understand how the cloud came to be, and how it relates to the accretion and feedback story of the Milky Way.
16
7
1607.06465
We present medium-resolution, near-ultraviolet Very Large Telescope/FLAMES observations of the star USNO-A0600-15865535. We adapt a standard method of stellar typing to our measurement of the shape of the Balmer ɛ absorption line to demonstrate that USNO-A0600-15865535 is a blue horizontal branch star, residing in the lower stellar halo at a distance of 4.4 kpc from the Sun. We measure the H &amp; K lines of singly ionized calcium and find two isolated velocity components, one originating in the disk, and one associated with the high-velocity cloud complex WD. This detection demonstrated that complex WD is closer than ∼4.4 kpc and is the first distance constraint on the +100 km s<SUP>-1</SUP> Galactic complex of clouds. We find that complex WD is not in corotation with the Galactic disk, which has been assumed for decades. We examine a number of scenarios and find that the most likely scenario is that complex WD was ejected from the solar neighborhood and is only a few kiloparsecs from the Sun.
false
[ "complex WD", "Sun", "clouds", "stellar typing", "∼4.4 kpc", "the high-velocity cloud complex WD", "decades", "USNO", "a blue horizontal branch star", "K lines", "the first distance constraint", "first", "Galactic", "USNO-A0600", "corotation", "Balmer", "two isolated velocity components", "the Sun", "scenarios", "the Galactic disk" ]
12.168644
9.088575
176
12408832
[ "Evslin, Jarah" ]
2016ApJ...826L..23E
[ "A Robust Measure of Dark Matter Halo Ellipticities" ]
2
[ "Institute of Modern Physics, CAS, NanChangLu 509, Lanzhou 730000, China" ]
[ "2017ApJ...841...90E", "2024A&A...683A.254D" ]
[ "astronomy" ]
3
[ "astrometry", "Local Group", "Astrophysics - Astrophysics of Galaxies" ]
[ "1987MNRAS.224...13D", "2002ApJ...564...60M", "2007PhRvD..75h3526S", "2009AJ....137.3100W", "2010MNRAS.402.1126T", "2011A&A...528A.119D", "2011MNRAS.416.1377V", "2012AJ....144....4M", "2012ApJ...755..145H", "2012JCAP...05..030S", "2012MNRAS.425.2069M", "2013MNRAS.430..105P", "2014EAS....67...23D", "2014MNRAS.439.2863V", "2015MNRAS.452L..41E", "2015MNRAS.454.3996D", "2015RAA....15.1945S" ]
[ "10.3847/2041-8205/826/2/L23", "10.48550/arXiv.1607.03595" ]
1607
1607.03595_arXiv.txt
\label{sec:intro} Cold dark matter simulations agree that dark matter halos \citep{triaxaq,triaxmil} and their subhalos \citep{darksub} are not spherical, but instead constant density surfaces are roughly triaxial ellipsoids whose axes vary from surface to surface. These shapes may be probed most cleanly in dwarf spheroidal (dSph) galaxies, the nearest of which are hosted by subhalos of the Milky Way's dark matter halo, as these often consist of 99\% dark matter or more. Many dSphs have multiple stellar populations which themselves are triaxial, with axes which are distinct from each other \citep{fornax2ell} and therefore likely from that of the dark matter halo. On the other hand, more exotic dark matter models, including many self-interacting dark matter models and Bose Einstein Condensate models, often have more spherical dark matter halo profiles \citep{sfer}, although quantifying this effect can be difficult \citep{kapsfer}. Therefore a measurement of the asphericity would provide a critical test of the cold dark matter paradigm. So far, the most robust determinations of halo shapes have been obtained for galaxy clusters. Such determinations are quite complicated, requiring a combination of X-ray, strong lensing and Sunyaev-Zeldovich effect observations, as in \citet{triaxcluster}. There is also a large literature which attempts to determine the shape of the Milky Way's own dark matter halo. However different methods have led to differing results, as is reviewed for example in \citet{triaxaq}. In the clean case of the dark matter subhalos inhabited by Milky Way satellites, the state of the art subhalo shape determination is \citet{triaxhayashi}. In this paper the authors found that the halo shapes are far more aspherical than those in CDM simulations, thus suggesting strong tension with the CDM paradigm. However, this paper makes several very strong assumptions, including an alignment of the stellar and dark matter halos, that the halos have two equal semi-principal axes and, even more critically, an assumption on the anisotropy of the stellar velocities which, in the spherical case, reduces to the assumption that the stellar velocities are isotropic. A shift in the stellar velocity anisotropy has, as the authors themselves described, a similar effect on the observables as a shift in the halo ellipticity and so this assumption alone may be responsible for the conclusions of the study. Unfortunately, the assumption cannot be lifted because the anisotropy itself cannot be determined using the line of sight stellar motions together with the Jeans equations \citep{degen}. However this situation is about to change. Right now the Gaia space telescope is measuring not only the line of sight velocities of many stars in each of these satellite galaxies, but for the first time it is also measuring their transverse velocities. With the full three-dimensional velocities in hand, this degeneracy will be broken and so more robust determinations of the halo shapes may be attempted. The precisions of these proper motion determinations will improve dramatically with the completion of the Thirty Meter Telescope (TMT) \citep{tmtdsc}. In particular, using the instrument IRIS, with a single view of such a galaxy one can obtain 20 $\mu$as precision relative positions of all of the stars bright enough for Gaia's astrometry. Combining this with Gaia's observations of the relative positions, for example six years earlier, one obtains an improvement in the velocity measurements of roughly a factor of four \citep{megaia}. With a second epoch of TMT observations several years later, one to two orders of magnitude more stars become available and the astrometric precision drops well below the stellar dispersions, which means that the measurements are limited primarily by the stellar dispersion and not the measurement error. In summary, our knowledge of stellar proper motions will have three great leaps. First, in two or three years when Gaia astrometry results become available, then when these results are combined with the first epoch of TMT observations and finally when the second epoch of TMT observations are combined with the first. In the rest of this letter we will show that Gaia will be unable to meaningfully constrain the halo shapes, however the first epoch of TMT observations can distinguish spherical from triaxial CDM halos inhabited by Sculptor-like dwarf spheroidal (dSph) galaxies with up to $5\sigma$ of precision for a favorable orientation, although for a generic orientation the confidence is around 3$\sigma$. However, since there are several such systems available, these 3$\sigma$ hints can be combined into a robust signal. On the other hand, we will see that the second epoch of TMT observations will provide an overwhelming and robust signal which excludes, at the 5$\sigma$ level, either CDM or else spherical halos.
16
7
1607.03595
In simulations of the standard cosmological model (ΛCDM), dark matter halos are aspherical. However, so far the asphericity of an individual galaxy’s halo has never been robustly established. We use the Jeans equations to define a quantity that robustly characterizes a deviation from rotational symmetry. This quantity is essentially the gravitational torque and it roughly provides the ellipticity projected along the line of sight. We show that the Thirty Meter Telescope (TMT), with a single epoch of observations combined with those of the Gaia Space Telescope, can distinguish the ΛCDM value of the torque from zero for each Sculptor-like dwarf galaxy with a confidence between 0 and 5σ, depending on the orientation of each halo. With two epochs of observations, TMT will achieve a 5σ discovery of torque and thus asphericity for most such galaxies, thus providing a new and powerful test of the ΛCDM model.
false
[ "dark matter halos", "most such galaxies", "rotational symmetry", "torque", "sight", "5σ", "Meter Telescope", "ΛCDM", "TMT", "the ΛCDM model", "observations", "each halo", "an individual galaxy’s halo", "the ΛCDM value", "the Gaia Space Telescope", "the standard cosmological model", "a 5σ discovery", "the gravitational torque", "the Thirty Meter Telescope", "each Sculptor-like dwarf" ]
12.254133
4.252724
-1
4855650
[ "Nicastro, F.", "Senatore, F.", "Krongold, Y.", "Mathur, S.", "Elvis, M." ]
2016arXiv160708364N
[ "The Milky Way Hot Baryons and their Peculiar Density Distribution: a Relic of Nuclear Activity" ]
0
[ "-", "-", "-", "-", "-" ]
null
[ "astronomy" ]
4
[ "Astrophysics - Astrophysics of Galaxies" ]
[ "2003ApJ...590L..33L", "2005ApJ...619...60L", "2011MNRAS.415...11D", "2012MNRAS.419.2251M", "2012MNRAS.425..605F", "2013A&ARv..21...61R", "2013ApJ...762...20F", "2013ApJ...770..118M", "2016A&A...594A..13P", "2016ApJ...828L..12N", "2016MNRAS.457..676N", "2017ApJ...835...52F" ]
[ "10.48550/arXiv.1607.08364" ]
1607
1607.08364_arXiv.txt
The visible baryonic mass of the Milky Way amounts to M$_b^{Obs} \simeq 0.65×\times 10^{11}$ M$_{\odot}$ $^{1)}$. The total baryonic plus dark matter mass of our Galaxy, is instead M$_{Tot} \simeq (1-2) \times 10^{12}$ M$_{\odot}$ $^{(2)}$. This, assuming a universal baryon fraction of $f_b = 0.157$ $^{(3)}$, implies a total baryonic mass of M$_b^{Pred} \simeq (1.6-3.2) \times 10^{11}$ M$_{\odot}$, between 2.5 and 5 times larger than observed. A large fraction of the baryonic mass of our Galaxy is thus currently eluding detection. This “missing-baryon” problem is not a monopoly of the Milky Way: most of the galaxies in the local universe suffer a deficit of baryonic mass compared to their dynamical mass and the problem is more serious at smaller dynamical masses (e.g. (4)), suggesting that lighter galaxies fail to retain larger fractions of their baryons. These baryons could be at least partly hiding under the form of tenuous hot ($\sim 10^6$ K) gas, heated up by recurring episodes of nuclear activity during the galaxy's lifetime, such as bursts of star formation followed by powerful supernova explosions or accretion-powered ignitions of the central supermassive black hole, which may have powered energetic outflows that pushed material out to large distances from the Galaxy's center. Over the past several years, a number of experiments, as well as theoretical works, have attempted to gain insights into the location and mass of the hot medium in our own Galaxy. Our peripheral position in the Galaxy, at about 8.5 kpc from the Galaxy's center and roughly in the Galaxy's plane, gives us hope of solving the problem: once a physically motivated density profile is assumed for the hot absorbing medium, the observed column densities (as well as the other observables) will depend critically on the sky position (and distance, for Galactic background targets) of the sources towards which the column densities are measured. This consideration has recently motivated several studies, which have used available spectra of extragalactic targets, with no other selection criterion than being at high Galactic latitudes, to measure OVII column densities and compare them with physically motivated or simple phenomenological density profile models $^{(5,6,7)}$. The results, however, are often contradictory, with estimated total masses of the million degree medium within a 1.2 virial-radius sphere (300 kpc) that strongly depend on the flatness of the assumed density profile, and range from a negligible M$_{Hot} \simeq 2.4 \times 10^9$ M$_{\odot}$ $^{(5)}$ to a significant M$_{Hot} \simeq 10^{11}$ M$_{\odot}$ $^{(7)}$. \noindent Here we present an experiment that settles the controversy by adopting a number of novel and rigourous data analysis techniques and sample selection criteria, and that has recently been published by the ApJ Letter ($^{(8)}$, hereinafter N16). \noindent Throughout this contribution, we refer to all densities and masses in units of ($A_O/4.9\times 10^{-4})^{-1} \times [Z/(0.3 Z_{\odot})]^{-1} (f_{OVII}/0.5)^{-1}$, where $A_O$ is the relative abundance of oxygen compared to hydrogen, $Z$ is the metallicity and $f_{OVII}$ is the fraction of OVII relative to oxygen. For easy comparison to other works (e.g. $^{(7)}$), we compute hot baryon masses within a 1.2 virial-radius sphere. Errors on best-fitting parameters (and quantities derived from those) are provided at 90\% confidence level for a number of interesting parameters equal to (31-N$_{dof}$), where N$_{dof}$ is the number of degrees of freedom in the fit.
16
7
1607.08364
We know that our Galaxy is permeated by tenuous, hot, metal-rich gas. However much remains unknown about its origin, the portion of the Galaxy that it permeates, its total mass, as any role it may play in regulating activity in the Galaxy. In a Letter currently in the press with the ApJ, we show that this hot gas permeates both the disk of the Galaxy and a large spherical volume, centered on the Galactic nucleus, and extending out to distances of at least 60-200 kpc from the center. This gas displays a peculiar density distribution that peaks about 6 kpc from the Galaxy's center, likely witnessing a period of strong activity of the central super-massive black hole of the Milky Way that occurred 6 Myrs ago. With our study we are also able to update the total baryonic mass of the Galaxy to Mb = (0.8-4)x1e11 Solar Masses, sufficient to close the Galaxy's baryon census.
false
[ "Galaxy", "strong activity", "the central super-massive black hole", "activity", "Solar Masses", "distances", "the Galaxys baryon census", "our Galaxy", "the Galaxy", "the Galaxys center", "=", "the Milky Way", "Mb", "a large spherical volume", "the total baryonic mass", "a peculiar density distribution", "the Galactic nucleus", "tenuous, hot, metal-rich gas", "the center", "This gas" ]
12.689284
8.107778
-1
12594004
[ "Sudoh, Takahiro", "Totani, Tomonori", "Makiya, Ryu", "Nagashima, Masahiro" ]
2017MNRAS.464.1563S
[ "Testing anthropic reasoning for the cosmological constant with a realistic galaxy formation model" ]
9
[ "Department of Astronomy, School of Science, The University of Tokyo, Hongo, Tokyo 113-0033, Japan", "Department of Astronomy, School of Science, The University of Tokyo, Hongo, Tokyo 113-0033, Japan", "Kavli Institute for the Physics and Mathematics of the Universe (Kavli IPMU, WPI), Todai Institutes for Advanced Study, The University of Tokyo, Kashiwa 277-8583, Japan; Max-Planck-Institut für Astrophysik, Karl-Schwarzschild Str. 1, D-85741 Garching, Germany", "Faculty of Education, Bunkyo University, Koshigaya, Saitama 343-8511, Japan" ]
[ "2017PhRvD..96h4062T", "2018MNRAS.477.3727B", "2018MNRAS.477.3744S", "2019AN....340..173T", "2019AsBio..19..126T", "2019PhR...807....1A", "2021MNRAS.508.5802S", "2022MNRAS.517...59O", "2022arXiv220908783K" ]
[ "astronomy", "physics" ]
7
[ "galaxies: formation", "cosmological parameters", "cosmology: theory", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]
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[ "10.1093/mnras/stw2401", "10.48550/arXiv.1607.00180" ]
1607
1607.00180_arXiv.txt
\label{sec:intro} The early observational indications for non-vanishing cosmological constant, $\Lambda$ \citep{Efstathiou1990,Fukugita1990,Yoshii1993,Krauss1995,Ostriker1995}, were further strengthened by type Ia supernova observations \citep{Riess1998,Perlmutter1999}, and established by the WMAP data of the cosmic microwave background radiation \citep{Spergel2003}, leading to the standard $\Lambda$CDM model that are consistent with many further high-precision observational tests until now \citep[see e.g. ][for reviews]{Frieman2008,Weinberg2013}. The cosmological constant is interpreted as the vacuum energy density, but theoretically natural values expected by particle physics are larger than the observed value $\Lambda_{\rm obs}$ by a factor of at least $\sim 10^{55}$ \citep[see e.g. ][for theoretical reviews]{Weinberg1989,Carroll2001,Sahni2002,Caldwell2009}. The discrepancy becomes even much worse if we assume that the natural value of the cosmological constant is set by the Planck energy density, which is $10^{123}$ times larger than the energy density corresponding to $\Lambda_{\rm obs}$. This is the so-called smallness problem. Furthermore, $\Lambda$ is not exactly zero but has a finite value, and we are living in a very special epoch when the energy densities of matter and vacuum become comparable in the long history of the universe. This is the coincidence problem. Many approaches have been proposed, ranging from a new field called dark energy to modification of the gravity theory, but there is no satisfactory solution yet for this problem. One approach to this problem is using the anthropic principle \citep[][see also \citealt{Banks1985}]{Barrow1986,Weinberg1987}. If $\Lambda \ll - \Lambda_{\rm obs}$, the universe would have collapsed much earlier than the present epoch \citep{Barrow1986}. On the other hand, if $\Lambda$ is positively too large, gravitational condensation of matter is suppressed by an accelerated expansion of the background universe, and hence no galaxy or intelligent life is formed \citep{Weinberg1987}. Such an idea is supported by some theories about very early universe suggesting a possibility that there is an ensemble of many multiverses and $\Lambda$ varies among different multiverses. It is reasonable to expect a nearly flat prior probability distribution $p(\Lambda)$ per unit $\Lambda$, because $\Lambda$ of a habitable universe must be smaller than the theoretically natural scale $\Lambda_{\rm th}$ by many orders of magnitude. If $p(\Lambda)$ is non-zero at $\Lambda = 0$, and $dp/d\Lambda \sim \Lambda_{\rm th}^{-1}$, $p(\Lambda)$ would be essentially constant within the range of $\Lambda$ for a habitable universe. Then not only the smallness but also the coincidence problem is solved because the probability of observing an absolute value smaller than $|\Lambda|$ scales as $\propto |\Lambda|$. The anthropic argument may not simply work if not only $\Lambda$ but also other physical constants or quantities are varied in different multiverses \citep{Tegmark1998,Aguirre2001,Graesser2004}, but in this work we consider the simplest scenario that only $\Lambda$ changes in a flat universe. Recently, it has been shown that only $\Lambda$ changes with a flat $p(\Lambda)$ among different homogeneous patches of the universe created by inflation, if the theory of gravity is extended to allow any inhomogeneous initial conditions about spacetime and matter, in contrast to general relativity in which the four constraints in the Einstein field equations limit the possible initial conditions \citep{Totani2016}. If we know the expected number density of observers in a universe with the cosmological constant $\Lambda$, $n(\Lambda)$, we can calculate the probability distribution $P(\Lambda)$ per unit $\Lambda$ for an observer realized in a universe, as \begin{equation} P(\Lambda)=\frac{n(\Lambda) p(\Lambda)}{\int_{0}^\infty n(\Lambda) p(\Lambda)d\Lambda} \ , \label{eq1} \end{equation} where $n(\Lambda)$ is a comoving density scaled to an early epoch when the effect of $\Lambda$ is negligible, to correct the difference of late expansion factor by changing $\Lambda$. Here we consider only $\Lambda \ge 0$, because the extension of $n(\Lambda)$ to the negative range is dependent on the formation time scale of an intelligent observer, which is highly uncertain. If it takes $\sim 5$ Gyr (appearance of human being after the formation of the Earth), $n(\Lambda)$ should rapidly drop below $\Lambda \lesssim - \Lambda_{\rm obs}$, and hence ignoring the rather narrow range of $-\Lambda_{\rm obs} \lesssim \Lambda < 0$ does not seriously affect the probability calculation \citep[][see also \citealt{Peacock2007}]{Weinberg1996}. The probability to observe $\Lambda$ smaller than the observed value is then \begin{equation} P(< \Lambda_{\rm obs}) \equiv \int_0^{\Lambda_{\rm obs}} P(\Lambda) d\Lambda \ . \end{equation} If this is not small, the anthropic principle can be a viable solution to the cosmological constant problem. Previous studies indeed show that $P(<$$\Lambda_{\rm obs})$ is not extremely small (typically 1--10\%), based on the structure formation theory of the universe with cold dark matter (CDM) \citep{Efstathiou1995,Martel1998,Garriga2000,Peacock2007}. In these studies the total dark matter mass included in collapsed dark haloes was used as an estimator of $n(\Lambda)$, and it was calculated analytically by formulations like the Press-Schechter theory, but baryonic physics related to formation of galaxies (e.g., gas cooling, star formation, supernova feedback, starbursts by galaxy mergers, metal production and chemical evolution) were not considered, though it should also be important to estimate $n(\Lambda)$. Furthermore, these studies assumed a minimum mass threshold for dark haloes that can harbor life, with a value similar to that of our Galaxy, to get a probability $P(<$$\Lambda_{\rm obs})$ that is not extremely small. However, such a treatment is clearly ad hoc, and we know that the mass distribution of galaxies extends to dwarf galaxies that are smaller than our Galaxy by a factor of 1000. In this work, we calculate $n(\Lambda)$ in a wide range of $\Lambda$ by using a semi-analytic model of galaxy formation in the framework of cosmological structure formation in a CDM universe. In such a model, the baryonic processes of galaxy formation mentioned above are taken into account, and model parameters are determined to reproduce a variety of observed data including luminosity functions, galaxy number counts, and several empirical relationships like the Tully-Fisher relation \citep[see e.g.][for a review]{Baugh2006}. The aim of this work is to examine whether a reasonably large $P(<$$\Lambda_{\rm obs})$ is obtained without an ad hoc galaxy mass threshold, when we take into account realistic physical processes of galaxy formation. It is known that the faint-end slope of galaxy luminosity function is flatter than that of dark haloes, and supernova feedback is believed as the primary mechanism working preferentially in low mass haloes to reduce the amount of stars. This would have a similar effect to the mass threshold in the previous studies, but a quantitative computation is necessary, which is the distinctive feature of this work. We first calculate $n(\Lambda)$ assuming that the number of life systems in a galaxy, $N_{\rm life}$, is proportional to stellar mass of the galaxy, $M_*$. However, formation of a star is not a sufficient condition of a habitable system for an observer. Clearly, we do not expect formation of a terrestrial planet or life around a zero metallicity star. Then there should be a metallicity dependence for the probability of finding an observer in a stellar system. It is known that massive galaxies generally have high metallicities, and this trend is reproduced by galaxy formation models. Then high-metallicity preference of life formation may also work as an effective threshold in galaxy mass, and hence help the anthropic argument of $\Lambda$. Metallicity evolution of galaxies is calculated in the galaxy formation model used in this work, and we will also study this effect quantitatively. The outline of this paper is as follows. In Section \ref{sec:methods}, we describe the galaxy formation model that we use, and methods of calculating $n(\Lambda)$ in universes with different $\Lambda$. Section \ref{sec:results} presents our main results, followed by a summary in Section \ref{sec:conclusion}. We adopt the "Planck+WP'' values reported in Table 2 of \citet{Planck} for the cosmological parameters of the present universe that we observe.
\label{sec:conclusion} In this work we examined the anthropic argument to explain the cosmological constant problem by a semi-analytic model of cosmological galaxy formation, assuming that an observable universe is created with a variable value of $\Lambda$ obeying to a nearly flat prior probability distribution per unit $\Lambda$, while any other physical parameter does not change. The galaxy formation model used here produces a mock catalog of dark matter haloes and their merger history by Monte-Carlo simulations based on the structure formation theory in the $\Lambda$CDM universe. Various astrophysical processes such as gas cooling, star formation, and supernova feedback are phenomenologicaly modeled to produce galaxies in dark haloes with physical quantities such as stellar mass and metallicity. Astrophysical model parameters have been determined to reproduce various observed data, and they are assumed not to change for different $\Lambda$. Assuming that the number of observers $N_{\rm life}$ is proportional to stellar mass $M_*$ in a galaxy, we find a median in the probability distribution $P(\Lambda)$ to be $\Lambda/\Lambda_{\rm obs} = 11.0$, and the probability of finding $\Lambda \le \Lambda_{\rm obs}$ to be $P(<\Lambda_{\rm obs})$ = 6.7\%. It should be noted that we obtained this result without introducing any galaxy mass threshold, which is in contrast to previous results based only on the formation history of dark matter haloes. Using the PS formalism and assuming that the number of observers is proportional to dark matter mass of collapsed haloes, we confirmed the previous results that a mass threshold close to our Galaxy halo must be assumed to get a probability that is not extremely small: $P(<\Lambda_{\rm obs})$ = 3.4\% for $M_{\rm DM, th} = 5 \times 10^{11} M_\odot$, though there exist much smaller galaxies. If we take a smaller threshold of $5 \times 10^8 M_\odot$, the probability reduces to 0.58\%. Our result using the galaxy formation model can be understood by the supernova feedback taken into account in the model; a significant fraction of dark matter mass is distributed in small mass haloes, but star formation in such haloes is suppressed by the feedback. We also tested the possibility that the number of observers depends on metallicity of galaxies. Introducing a metallicity threshold does not change the probability $P(<\Lambda_{\rm obs})$ significantly; it increases from 7.3 to 9.0\% for $Z_{\rm th} = $ 0.1 to 1.0 $Z_\odot$. This is because low metallicity galaxies are generally small galaxies (the mass-metallicity relation), and such galaxies do not include a significant fraction of stellar mass in the universe due to the supernova feedback. If we assume that the number of observers is proportional to $N_{\rm life} \propto M_*Z$ of galaxies, we found $P(<\Lambda_{\rm obs})$ = 9.7\%. We conclude that a reasonable estimate of the probability to find a small $\Lambda$ as observed is 7--10\%, which is not extremely small, based on a realistic model of galaxy formation. Therefore the anthropic argument is a viable explanation for the cosmological constant problem. If, in future, a convincing theory is established by fundamental physics predicting that only $\Lambda$ is variable with a flat prior distribution when a universe is created, the anthropic argument may become the leading candidate for the solution of the cosmological constant problem.
16
7
1607.00180
The anthropic principle is one of the possible explanations for the cosmological constant (Λ) problem. In previous studies, a dark halo mass threshold comparable with our Galaxy must be assumed in galaxy formation to get a reasonably large probability of finding the observed small value, P(&lt;Λ<SUB>obs</SUB>), though stars are found in much smaller galaxies as well. Here we examine the anthropic argument by using a semi-analytic model of cosmological galaxy formation, which can reproduce many observations such as galaxy luminosity functions. We calculate the probability distribution of Λ by running the model code for a wide range of Λ, while other cosmological parameters and model parameters for baryonic processes of galaxy formation are kept constant. Assuming that the prior probability distribution is flat per unit Λ, and that the number of observers is proportional to stellar mass, we find P(&lt;Λ<SUB>obs</SUB>) = 6.7 per cent without introducing any galaxy mass threshold. We also investigate the effect of metallicity; we find P(&lt;Λ<SUB>obs</SUB>) = 9.0 per cent if observers exist only in galaxies whose metallicity is higher than the solar abundance. If the number of observers is proportional to metallicity, we find P(&lt;Λ<SUB>obs</SUB>) = 9.7 per cent. Since these probabilities are not extremely small, we conclude that the anthropic argument is a viable explanation, if the value of Λ observed in our Universe is determined by a probability distribution.
false
[ "cosmological galaxy formation", "galaxy formation", "galaxy luminosity functions", "galaxies", "SUB", "other cosmological parameters", "model parameters", "P(&lt;Λ", "Λ", "stars", "obs</SUB", "much smaller galaxies", "any galaxy mass threshold", "many observations", "baryonic processes", "the observed small value", "stellar mass", "metallicity", "a semi-analytic model", "observers" ]
10.926347
0.474067
89
12409073
[ "Pavesi, Riccardo", "Riechers, Dominik A.", "Capak, Peter L.", "Carilli, Christopher L.", "Sharon, Chelsea E.", "Stacey, Gordon J.", "Karim, Alexander", "Scoville, Nicholas Z.", "Smolčić, Vernesa" ]
2016ApJ...832..151P
[ "ALMA Reveals Weak [N II] Emission in \"Typical\" Galaxies and Intense Starbursts at z = 5-6" ]
71
[ "Department of Astronomy, Cornell University, Space Sciences Building, Ithaca, NY 14853, USA ;;", "Department of Astronomy, Cornell University, Space Sciences Building, Ithaca, NY 14853, USA", "Spitzer Science Center, California Institute of Technology, MC 220-6, 1200 East California Boulevard, Pasadena, CA 91125, USA", "National Radio Astronomy Observatory, PO Box O, Socorro, NM 87801, USA ; Cavendish Astrophysics Group, University of Cambridge, Cambridge, CB3 0HE, UK", "Department of Astronomy, Cornell University, Space Sciences Building, Ithaca, NY 14853, USA", "Department of Astronomy, Cornell University, Space Sciences Building, Ithaca, NY 14853, USA", "Argelander-Institut für Astronomie, Universität Bonn, Auf dem Hügel 71, D-53121 Bonn, Germany", "Astronomy Department, California Institute of Technology, MC 249-17, 1200 East California Boulevard, Pasadena, CA 91125, USA", "University of Zagreb, Physics Department, Bijenička cesta 32, 10002 Zagreb, Croatia" ]
[ "2016ApJ...829L..11P", "2017A&A...604A..53C", "2017A&A...605A..42C", "2017ApJ...834L..16U", "2017ApJ...842L..16L", "2017ApJ...845...41B", "2017ApJ...845..154V", "2017ApJ...847...21F", "2017ApJ...850..180J", "2017MNRAS.469L..16L", "2018ApJ...861...43P", "2018ApJ...861...95H", "2018MNRAS.474.1718N", "2018MNRAS.481...59Z", "2018MNRAS.481.1631B", "2018Natur.553..178S", "2019A&A...631A.167D", "2019ApJ...871...85L", "2019ApJ...872..117G", "2019ApJ...875....6G", "2019ApJ...876....1T", "2019ApJ...882..168P", "2019ApJ...883L..29L", "2019ApJ...886...29S", "2019ApJ...886...60S", "2019MNRAS.488.1489H", "2020A&A...637A..32B", "2020A&A...643A...1L", "2020ApJ...895...81R", "2020ApJ...896...93H", "2020ApJ...898...33C", "2020ApJ...900..131L", "2020ApJ...902...37D", "2020ApJ...902..112B", "2020ApJ...904..131N", "2020ApJS..247...61F", "2020IAUS..341..231H", "2020IAUS..341..283L", "2020MNRAS.494.4090C", "2020MNRAS.496..875R", "2020RSOS....700556H", "2021A&A...646A..68S", "2021A&A...646A..76L", "2021A&A...649A..31H", "2021A&A...649A.152K", "2021A&A...650C...2S", "2021A&A...652A..66P", "2021ApJ...907...62R", "2021ApJ...911...99F", "2021ApJ...912...73A", "2021ApJ...913...41L", "2021ApJ...916...12B", "2021ApJ...923....5S", "2021MNRAS.507.3540J", "2022A&A...664A..73B", "2022ApJ...925..174V", "2022ApJ...927..152M", "2022ApJ...928..179L", "2022MNRAS.510.5560S", "2022MNRAS.513.3122S", "2022MNRAS.515.1751W", "2022MNRAS.517.2508W", "2022Univ....8..314F", "2023A&A...673A.157D", "2023A&A...679A.131R", "2023pcsf.conf...30V", "2024ApJ...965..123L", "2024ApJ...968..113E", "2024arXiv240105487R", "2024arXiv240405352C", "2024arXiv240512955J" ]
[ "astronomy" ]
13
[ "cosmology: observations", "galaxies: formation", "galaxies: high-redshift", "galaxies: ISM", "galaxies: starburst", "radio lines: galaxies", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1985ApJ...291..722T", "1990ApJ...358..116W", "1994A&AS..103...97M", "1996ARA&A..34..279S", "1998ApJ...504L..11L", "1999ApJS..123....3L", "2003ApJ...588...65S", "2003ApJ...594..758L", "2004ApJ...604..125B", "2005A&A...440L..45W", "2005A&A...440L..51M", "2006ApJ...644..283K", "2006ApJ...652L.125O", "2006ApJ...653..977S", "2007ApJ...665..936C", "2007MNRAS.382..325B", "2009ApJ...698.1010B", "2009ApJ...703..785D", "2009MNRAS.394.1812P", "2009MNRAS.395.1476R", "2009Natur.457..451D", "2009Natur.457..699W", "2010ApJ...720L.131R", "2010ApJ...724L.153R", "2010MNRAS.405....2L", "2010MNRAS.407.1464C", "2011A&A...526A.149N", "2011ARA&A..49..525S", "2011ApJ...733...29S", "2011ApJ...733..101G", "2011ApJ...739..100O", "2011ApJ...739L..31R", "2011Natur.470..233C", "2012A&A...538L...4C", "2013A&A...554A.103P", "2013A&A...557A..95R", "2013ARA&A..51..105C", "2013ARA&A..51..207B", "2013ApJ...763..120C", "2013ApJ...765L..13Z", "2013ApJ...774...68D", "2013ApJ...774L..10K", "2013ApJ...776...38F", "2013ApJ...776L..24T", "2013ApJ...778....2S", "2013MNRAS.432..455D", "2013MNRAS.433.1567V", "2013Natur.496..329R", "2013RMxAA..49..137F", "2014A&A...565A..59D", "2014A&A...568A..62D", "2014ApJ...782L..17D", "2014ApJ...783...59R", "2014ApJ...795..165S", "2014ApJ...795..174K", "2014ApJ...796...84R", "2014ApJS..214...15S", "2014MNRAS.439.2096W", "2014PhR...541...45C", "2015A&A...578A..53C", "2015ApJ...799...21S", "2015ApJ...799..209W", "2015ApJ...801...72R", "2015ApJ...802L..11L", "2015ApJ...807..180W", "2015ApJ...814..133G", "2015Natur.522..455C", "2016A&A...586L...7B", "2016A&A...590A..43L", "2016A&A...591A..33R", "2016ApJ...819...69Z", "2016ApJ...821..122F", "2016ApJ...826..159S", "2016ApJ...826..175H", "2016ApJ...828...18M", "2016ApJ...829...93K", "2016ApJS..224...24L", "2016MNRAS.458.3466V", "2017MNRAS.467.1300V" ]
[ "10.3847/0004-637X/832/2/151", "10.48550/arXiv.1607.02520" ]
1607
1607.02520_arXiv.txt
The first billion years after the Big Bang is a crucial epoch for understanding galaxy evolution because we can directly witness the initial stages of galaxy assembly. In contrast to present day, most galaxies at these epochs are thought to have only formed a small fraction of their final stellar mass, to be accreting pristine gas very actively from the cosmic web, to have low metal abundances, and to have been affected in their star formation properties by the pristine quality of the inter-stellar medium (ISM) \citep[e.g.,][]{BH07,P09, L10, CN10, RW09, B09, S06}. Investigating the ISM and its relationship to star formation during the first billion years of cosmic time is a promising test bed for galaxy formation models \citep[e.g.,][]{D09a, D09b, D13}, complementing studies of the peak epoch of galaxy assembly ($z\sim2$--$3$; e.g., \citealt{Sh11,CW13,Cas14}). To faithfully model the interplay of the physical processes at the root of galaxy assembly and evolution we need observations of accurate diagnostics of the different physical phases of the gas. While luminous starbursting galaxies like submillimeter galaxies (SMGs) and quasar hosts have been targeted for more than a decade at $z>5$ (e.g., \citealt{M05,W09,R13}; see \citealt{CW13} for a review), we are only now reaching the capability of investigating the ISM in ``normal" star-forming galaxies that are more representative of the general galaxy population at these epochs, in particular Lyman-break galaxies (LBGs; \citealt{R14}, hereafter R14; \citealt{C15}, hereafter C15; \citealt{Wi15}). CO excitation ladders have become a routine tool to investigate the physical conditions in molecular gas for starbursts \citep[e.g.,][]{R10a,R11,R13,We05, Sc11}; however, the prospects for detection in normal galaxies at $z>5$ are not clear due to metallicity effects \citep[][R14]{T13}. In fact, CO detections of LBGs to date, even exploiting strong lensing, have been limited to $z \lesssim3$ \citep[e.g.,][]{B04,Cop07,R10b,S13}. Submillimeter fine-structure lines of the most abundant atomic metal species (mainly C, N, O) and their ions offer a unique angle and an unobscured view of the ISM properties and conditions for star formation that is accessible to ALMA at high redshift. The far-infrared (FIR) cooling lines, less affected by dust attenuation than optical lines, are powerful probes of the star formation activity, linking them directly to the surrounding medium from which stars are born. Different lines, better in combination, can be used as diagnostics of the far-ultraviolet (FUV) flux, gas density, temperature, and filling factor of the photon-dominated regions (PDRs) and ionized regions \citep[e.g.,][]{TH85,W90,K06}. Recent surveys with the {\it Herschel Space Observatory} provide the context for high redshift studies by giving benchmark sets of local galaxies, for which almost complete suites of FIR lines yield solid constraints and allow detailed ISM modeling \citep[e.g.,][]{Ro15,Co15,DL14,S15,K16}. The ionization potential of nitrogen ($14.5\,$eV) is greater than that of hydrogen, therefore the singly ionized nitrogen lines, [N{\sc ii}], probe the effect of UV photons emitted by massive young stars. Since N$^+$ is only present in the ionized medium it is a tracer of the extended low density envelope of H{\sc ii} regions, the ionized surfaces of dense atomic and molecular clouds and the warm ionized medium (WIM). \cite{GS15} find that the dominant source of [N{\sc ii}] emission in the Milky Way is not the WIM, because the electron density they measure from [N{\sc ii}] fine-structure line ratios is two orders of magnitude higher than what is expected in the Galactic diffuse medium, but it is expected that different types of galaxies, at different ages, will differ significantly in their ISM phase structure. The high brightness of the [C{\sc ii}] $158\,\mu$m line makes it an ideal target at high redshift, where it allows dynamical studies and a probe into the ISM properties and star formation. However, [C{\sc ii}] can originate from a range of gas phases and the line luminosity alone gives no direct information on its origin. In fact, low-$z$ observations have shown that while most of the [C{\sc ii}] luminosity comes from the atomic PDRs and diffuse cold neutral medium (CNM), significant fractions also come from ionized gas regions and CO-dark molecular clouds \citep{P13}. \cite{O06,O11} suggested that the [N{\sc ii}] $205\,\mu$m line could be used in conjunction with [C{\sc ii}] $158\,\mu$m to separate the ionized from neutral fraction of the ISM, because the latter is emitted by both weakly ionized and atomic gas. The transition critical densities of these two lines in ionized gas are very similar, implying that the line ratio for the ionized medium is nearly constant ($\sim3$--4, \citealt{O06}) and set by the relative carbon to nitrogen abundance, making [N{\sc ii}] a tracer of the ionized fraction of [C{\sc ii}]-emitting gas. \cite{L16} combined [N{\sc ii}] and [C{\sc ii}] Galactic observations to determine the fraction of ionized gas contributing to the [C{\sc ii}] emission in the Milky Way for many lines of sight, finding this to vary mostly between 0.5 and 0.8 (corresponding to [C{\sc ii}]/[N{\sc ii}] ratios $\sim$4--10), where the Galactic center line of sights have the smallest ionized fraction, probably due to higher gas densities. The [N{\sc ii}] $205\,\mu$m line is the tool of choice to characterize the low-density ionized ISM at high redshift. Thus, we here use ALMA to measure the [N{\sc ii}] properties of a sample spanning the variety of star-forming activity known at $z=5$--6: a hyper-luminous nuclear starburst (AzTEC-3; SFR$\,{\rm\sim1100\,M_{\odot}\,yr^{-1}}$), a dusty, high star formation rate LBG (HZ10; SFR$\,{\rm\sim170\,M_{\odot}\,yr^{-1}}$) and a typical, less star-forming LBG (LBG-1; SFR$\,{\rm\sim10-30\,M_{\odot}\,yr^{-1}}$). These last two galaxies are near $L^*_{\rm UV}$ for LBGs at $z=5$--6, and can be considered ``typical" because they are consistent with the current constraints on the star-forming main sequence of galaxies at these redshifts \citep[e.g.,][]{Sp14}. All galaxies were previously detected in [C{\sc ii}] with ALMA (R14; C15). The structure of this paper is as follows. In Section~\ref{obs} we describe the ALMA observations of the [N{\sc ii}] and [C{\sc ii}] lines utilized in this study. Our results are presented in Section~\ref{res}, and in Section~\ref{model} we describe the H{\sc ii} and PDR models we have utilized to assist our interpretation. In Section~\ref{analysis}, we present our interpretation for each galaxy in our sample. We close with a discussion of the impact of our findings and of this kind of FIR line studies in the first giga-year of cosmic time in Section~\ref{discuss}. In this work we assume a flat $\Lambda$CDM cosmology with $H_0$=$70\,$km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm m}$=0.3 and $\Omega_\Lambda$=0.7.
\label{discuss} We have detected [N{\sc ii}] $205\,\mu$m emission towards the highest redshift sample of galaxies to date. Our sample spans almost two orders of magnitude in star formation rate and includes the compact starbursting SMG AzTEC-3 ($z=5.3$; SFR$\,{\rm\sim1100\,M_{\odot}\,yr^{-1}}$; R14). From the combined R14 \& C15 sample of [C{\sc ii}]-detected LBGs we also selected the relatively dusty, above average SFR galaxy, HZ10 ($z=5.7$; SFR$\,{\rm\sim170\,M_{\odot}\,yr^{-1}}$), which lies at the upper limit, although within the scatter of the $z=5$--6 star-forming Main Sequence, and a typical lower-SFR galaxy, LBG-1 (SFR$\,{\rm\sim10-30\,M_{\odot}\,yr^{-1}}$), which we consider representative of the ``normal" population of star-forming galaxies at $z\sim5.3$ because it appears to lie on the Main Sequence. Our observations of the [C{\sc ii}]/[N{\sc ii}] luminosity ratio are summarized in Fig.~\ref{ratio_plot}, and are shown together with all other high-redshift measurements to date, as well as some local galaxy samples. We have re-measured the [N{\sc ii}] line luminosity in the BR1202--0725 system members (a QSO, an SMG and two LAEs at $z=4.7$) based on higher sensitivity archival ALMA data (project ID: 2013.1.00745.S) to update the ratios presented in \cite{D14}. The [N{\sc ii}] line is detected with high significance for the QSO, SMG, and LAE-2, but is not detected in LAE-1.\footnote{The [N{\sc ii}] line fluxes are $1.5 \pm 0.2 \,$Jy km s$^{-1}$ for the SMG, $0.74 \pm 0.07\,$Jy km s$^{-1}$ for the QSO, $<0.5\,$mJy peak ($3\sigma$ limit) for LAE-1, and $0.30 \pm 0.06\,$Jy km s$^{-1}$ for LAE-2.} The non-detection of [N{\sc ii}] emission in LAE-1 may not be strongly constraining if it is similar to LBG-1, with the caveat that it is likely to be affected by tidal forces from the SMG-QSO merger and the intense QSO radiation. Furthermore, the [N{\sc ii}] line in LAE-2 is very broad ($\sim1000\,$km s$^{-1}$), suggesting that the [C{\sc ii}] line measured in \cite{C13} is truncated due to the band edge, and thus may only cover $\sim$30\% of the full line emission.\footnote{Adopting the [C{\sc ii}] line luminosities from \cite{C13}, the [C{\sc ii}]/[N{\sc ii}] line ratios are $11\pm2$ for the SMG, $12\pm2$ for the QSO, $>13.5$ for LAE-1, and, when using the velocity range corresponding to the [C{\sc ii}] measurement, $11\pm2$ for LAE-2.} The very broad [N{\sc ii}] line in LAE-2 is compatible with the Ly$\alpha$ FWHM of $1225\pm257\,$km s$^{-1}$ in \cite{W14}, which may imply that the ISM in this galaxy is strongly affected by quasar winds. Recent work \citep[e.g.,][]{K13,S16} suggests that stellar metallicity in early galaxies might be very low compared to gas-phase metallicity (which is dominated by oxygen, nitrogen, and carbon abundances). This is due to the expected enrichment time for iron of the order of $1\,$Gyr; the iron abundance is responsible for the biggest contribution to stellar opacity and is not enriched in the ISM by core-collapse supernovae and is therefore expected to be low in galaxies at this epoch, hence greatly affecting massive star structure and UV radiative output. Since we expect the gas properties of the ISM under investigation to depend on the hardness and intensity of the stellar radiation field, these properties may not be directly comparable to local galaxies. It is also expected that the high gas fractions typical for galaxies at these epochs might strongly modify the gas phase structure \citep[e.g.,][]{V13}. Figure~\ref{ratio_plot} shows that there seems to be a relatively narrow range (0.5 dex) for [C{\sc ii}]/[N{\sc ii}] ratios in SMGs and quasar-hosts at high-redshift. Their range is similar to, although perhaps slightly higher than, the ratios measured in local LIRGs and ULIRGs. This does not fit our expectations based on local trends with density and star formation rate surface density \citep[e.g.,][]{L15}; these trends would suggest a significantly higher ratio than we observe. We therefore suggest that the global measurements at high-redshift may not be observing the [N{\sc ii}] emission coming from the equivalent regions in local starbursts, where the H{\sc ii} regions at the site of recent star-formation dominate, but a more diffuse ISM component which might not be as prevalent in less gas-rich local analogs. If this were to be correct, it is possible that the ionized gas in the nuclear regions, all of which is surrounding the recently formed massive stars may be in a state of higher ionization. In this case both nitrogen and oxygen may predominantly be in their doubly ionized state as perhaps observed in LBG-1. It is likely that in order to unravel the physical origin of the line ratio under consideration, spatially resolved line studies are needed; spatial global averages might give a very biased view of the physical gas conditions if the neutral and the ionized gas are not co-spatial. HZ10 is the first high redshift galaxy where this type of study seems possible. In fact, we find tentative evidence for the [N{\sc ii}] emission originating from only part of the [C{\sc ii}]-emitting region, which is reminiscent of clumpy star-forming gas disks observed in other tracers of ionized gas like H$\alpha$ \citep[e.g.,][]{G11,W15}. Since rest-frame optical wavelength tracers like H$\alpha$ are presently not accessible in the first giga-year of cosmic time, FIR fine-structure lines like [N{\sc ii}] may become the tracer of choice for this kind of resolved clump/disk dynamic studies in the earliest galaxies. Due to the lower density and dust content relative to a compact starburst like AzTEC-3, regions of different ionization might be more easily accessible in less extreme $z>5$ galaxies, like HZ10. Future studies are necessary to determine if the regions with lower [N{\sc ii}] emission are dominated by neutral gas, perhaps at higher density than the [N{\sc ii}]-emitting fraction (as potentially observed in AzTEC-3), or whether [N{\sc iii}] dominates the nitrogen emission in these regions, suggesting LBG-1-like conditions. A different effect may become dominant for ``normal" galaxies at $z=5$--6, like LBG-1, where the low dust content of the ISM and the very young stellar population perhaps start affecting the phase structure of the low-metallicity gas. In fact, our data constrain the [C{\sc ii}]/[N{\sc ii}] ratio to be larger than $\sim 40$ and probably in the 60--100 range, indicative of a minor contribution from ionized gas to [C{\sc ii}] compared to that in local normal star-forming galaxies, for which the measured ratio is typically 10--25 \citep{Z16,DS}. Understanding the [C{\sc ii}] emission line in such galaxies appears crucial as the low FIR and CO emissions (R14, C15) suggest that [C{\sc ii}] might be an essential neutral gas coolant in low dust environments at $z>5$, which is necessary to achieve dense gas formation and hence star-formation. Our measurement is compatible with the recent findings of high ratios in local dwarfs \citep{Co15}, strengthening the case that these might be appropriate analogs for the ISM in early galaxies, in some respects. \cite{Co15} find [O{\sc iii}] to be the brightest FIR line in these galaxies, suggesting abundant ionized gas in a higher ionization state than what is traced by [N{\sc ii}]. Measuring either [O{\sc iii}] or [N{\sc iii}] will be necessary to determine the relative importance of the ionized ISM component in normal $z=5$--6 galaxies, and it could become an important tool due to its brightness, possibly comparable to [C{\sc ii}] or brighter \citep{I16}. The extremely low [N{\sc ii}] value measured in LBG-1 holds further evidence in favor of hard and intense radiation in early star-forming galaxies, with possibly much higher escape fraction of hard ionizing photons (dominating the ISM throughout the galaxy) due to low dust abundances, and therefore could potentially help to shed light on the mechanisms of cosmic re-ionization. \smallskip We thank the anonymous referee for a helpful and constructive report and T. K. Daisy Leung for helpful discussion. DR and RP acknowledge support from the National Science Foundation under grant number AST-1614213 to Cornell University. R.P. acknowledges support through award SOSPA3-008 from the NRAO. AK acknowledges support by the Collaborative Research Council 956, sub-project A1, funded by the Deutsche Forschungsgemeinschaft (DFG). DR acknowledges the hospitality at the Aspen Center for Physics and the Kavli Institute for Theoretical Physics during part of the writing of this manuscript. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. This paper makes use of the following ALMA data: ADS/JAO.ALMA\#2015.1.00928.S, 2011.0.00064.S, 2012.1.00523.S, 2011.0.00006.SV and 2013.1.00745.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada), NSC and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ
16
7
1607.02520
We report interferometric measurements of [N II] 205 μm fine-structure line emission from a representative sample of three galaxies at z = 5-6 using the Atacama Large (sub)Millimeter Array (ALMA). These galaxies were previously detected in [C II] and far-infrared continuum emission and span almost two orders of magnitude in star formation rate (SFR). Our results show at least two different regimes of ionized interstellar medium properties for galaxies in the first billion years of cosmic time, separated by their {L}<SUB>[{{C</SUB>}{{II}}]}/{L}<SUB>[{{N</SUB>}{{II}}]} ratio. We find extremely low [N II] emission compared to [C II] ({L}<SUB>[{{C</SUB>}{{II}}]}/{L}<SUB>[{{N</SUB>}{{II}}]}={68}<SUB>-28</SUB><SUP>+200</SUP>) from a “typical” ∼ {L}<SUB>{UV</SUB>}<SUP>* </SUP> star-forming galaxy, likely directly or indirectly (by its effect on the radiation field) related to low dust abundance and low metallicity. The infrared-luminous modestly star-forming Lyman-break galaxy (LBG) in our sample is characterized by an ionized-gas fraction ({L}<SUB>[{{C</SUB>}{{II}}]}/{L}<SUB>[{{N</SUB>}{{II}}]}≲ 20) typical of local star-forming galaxies and shows evidence for spatial variations in its ionized-gas fraction across an extended gas reservoir. The extreme SFR, warm and compact dusty starburst AzTEC-3 shows an ionized fraction higher than expected given its SFR surface density ({L}<SUB>[{{C</SUB>}{{II}}]}/{L}<SUB>[{{N</SUB>}{{II}}]}=22+/- 8) suggesting that [N II] dominantly traces a diffuse ionized medium rather than star-forming H II regions in this type of galaxy. This highest redshift sample of [N II] detections provides some of the first constraints on ionized and neutral gas modeling attempts and on the structure of the interstellar medium at z = 5-6 in “normal” galaxies and starbursts.
false
[ "galaxies", "galaxy", "C II", "ionized interstellar medium properties", "low dust abundance", "star formation rate", "low metallicity", "local star-forming galaxies", "star-forming H II regions", "SFR", "ALMA", "ionized and neutral gas modeling attempts", "cosmic time", "medium", "z", "spatial variations", "first", "[N II", "an extended gas reservoir", "Array" ]
13.333057
7.653903
132
12621336
[ "Derenzo, Stephen", "Essig, Rouven", "Massari, Andrea", "Soto, Adrián", "Yu, Tien-Tien" ]
2017PhRvD..96a6026D
[ "Direct detection of sub-GeV dark matter with scintillating targets" ]
149
[ "Lawrence Berkeley National Laboratory, Mail Stop 55-121, Berkeley, California 94720, USA", "Stony Brook University, Stony Brook, New York 11794, USA", "Stony Brook University, Stony Brook, New York 11794, USA", "Stony Brook University, Stony Brook, New York 11794-3800, USA; Stony Brook University, Stony Brook, New York 11794, USA", "Stony Brook University, Stony Brook, New York 11794, USA" ]
[ "2016EPJC...76..706J", "2016JInst..11P0003L", "2016arXiv160808632A", "2017IJMPA..3230023B", "2017JCAP...06..055D", "2017JCAP...10..031E", "2017JHEP...06..087B", "2017JHEP...08..078K", "2017JHEP...09..116B", "2017JHEP...10..013G", "2017JHEP...10..162C", "2017PhRvD..95b3013H", "2017PhRvD..95e6011E", "2017PhRvD..95e6019K", "2017PhRvD..95i5001B", "2017PhRvD..96c5009R", "2017PhRvD..96d3010M", "2017PhRvD..96d3017E", "2017PhRvD..96h3521C", "2017PhRvD..96i5022F", "2017PhRvD..96k5021K", "2017PhRvL.118c1803K", "2017arXiv170704591B", "2018IJMPA..3350135L", "2018JAP...123k4501D", "2018JHEP...09..153K", "2018JHEP...11..066G", "2018NIMPA.901....6C", "2018PhLB..785..386K", "2018PhRvD..97a5004H", "2018PhRvD..97i5029E", "2018PhRvD..97k5026W", "2018PhRvD..98b3005C", "2018PhRvD..98k5034G", "2018PhRvL.120m1801F", "2018PhRvL.121j1801D", "2018PhRvX...8d1001A", "2018arXiv180308146W", "2018arXiv180509345H", "2019EPJC...79...99B", "2019ITNS...66.2333V", "2019JCAP...09..070E", "2019JPhG...46j5002V", "2019PhRvD..99a5005H", "2019PhRvD..99i5023D", "2019PhRvD..99k5009D", "2019PhRvD.100c5025C", "2019PhRvL.123o1802H", "2019PhRvR...1c3105E", "2019arXiv190105822P", "2019arXiv190409362V", "2019arXiv190607541E", "2019arXiv190704324L", "2019arXiv191000550H", "2020JCAP...01..004S", "2020JCAP...05..036A", "2020JCAP...06..056C", "2020JHEP...02..134D", "2020JHEP...03..036T", "2020JHEP...07..081H", "2020PhRvD.101e5004G", "2020PhRvD.101e6001B", "2020PhRvD.101g6014B", "2020PhRvD.101l3012K", "2020PhRvD.102a5017K", "2020PhRvD.102b3038L", "2020PhRvD.102c5014B", "2020PhRvD.102d3007L", "2020PhRvL.124b1801E", "2020PhRvL.124r1301D", "2020PhRvL.124t1801T", "2020PhRvR...2c3195C", "2020PhyOJ..13..172E", "2020arXiv200913534T", "2020arXiv200914302C", "2021ApPhL.118b2601F", "2021JCAP...12..048C", "2021JHEP...09..123M", "2021JHEP...11..198C", "2021NIMPA.98964957D", "2021PhRvD.103c5001D", "2021PhRvD.103f3022C", "2021PhRvD.104a5031K", "2021PhRvD.104c6011B", "2021PhRvD.104e6009L", "2021PhRvD.104g6013B", "2021PhRvD.104g6020B", "2021PhRvD.104h3023K", "2021PhRvD.104i5015G", "2021PhRvD.104i6001L", "2021PhRvL.126i1804G", "2021PhRvL.127k1301D", "2021PhRvL.127o1802H", "2021PhRvR...3a3069I", "2021PhRvR...3c3149C", "2021Physi...3..473P", "2021ScPP...10...30B", "2021arXiv210904473H", "2021arXiv211002985H", "2021arXiv211108576U", "2022JHEP...05..034M", "2022JHEP...05..191G", "2022JHEP...07..044F", "2022JLTP..209..441T", "2022NIMPA103966981B", "2022PhRvD.105a5001T", "2022PhRvD.105a5010C", "2022PhRvD.105e5023E", "2022PhRvD.105i5009L", "2022PhRvD.106a5024C", "2022PhRvD.106b3026B", "2022PhRvD.106d3004L", "2022PhRvD.106g5004G", "2022PhRvD.106i5041D", "2022PhRvD.106k2005H", "2022PhRvL.128q1301R", "2022PhRvL.128s1801H", "2022PhRvL.129v1301Z", "2022PhRvX..12a1009D", "2022RPPh...85e6201B", "2022RPPh...85f6901K", "2022arXiv220308297E", "2022arXiv220310089E", "2022arXiv220315978C", "2022arXiv220906854G", "2022arXiv220907426C", "2023ApPhL.122x3506L", "2023IJMPA..3850144W", "2023JCAP...03..052C", "2023JCAP...11..024B", "2023JCAP...12..009L", "2023JHEP...01..023B", "2023PDU....4001221M", "2023PhLB..84037853G", "2023PhLB..84137922C", "2023PhRvD.107c5035T", "2023PhRvD.107f3002H", "2023PhRvD.107g6015H", "2023PhRvD.107i5035B", "2023PhRvD.108c5035P", "2024JCAP...03..047B", "2024PhRvD.109e5009D", "2024PhRvD.109i2005B", "2024PhRvD.109i5015G", "2024PhRvL.132r1801B", "2024arXiv240111971L", "2024arXiv240206817C", "2024arXiv240315860A", "2024arXiv240415741H" ]
[ "astronomy", "physics" ]
6
[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Experiment" ]
[ "1949PhRv...75..796H", "1949PhRv...75.1611H", "1964SSCom...2..353C", "1968JAP....39.2029K", "1970PSSAR...3..735K", "1979PhRvL..43.1494H", "1982JChPh..76..843D", "1986PhRvD..33.3495D", "1987ApPhL..51..406P", "1989PhRvL..63.1719L", "1990NIMPA.289..406L", "1991PhRvB..43.4187L", "1992SeScT...7.1271D", "1994PhST...54....7D", "1995PhRvL..74.2623S", "1996PhRvL..77.3865P", "2001PhRvD..64d3502S", "2002NIMPA.480..494A", "2003ApPhL..83..791M", "2003ITNS...50..767M", "2004NuPhB.683..219B", "2004PhRvL..93t1803B", "2005NIMPA.537..357M", "2005NIMPA.538..657Z", "2005NIMPA.553..578M", "2005PhRvD..72j3508M", "2006MPLA...21..457B", "2006PhRvD..74e4034F", "2006PhRvL..96n1802B", "2007PhLB..651..374S", "2007PhRvD..75k5017F", "2007PhRvL..99i1301L", "2008ITNS...55.1062M", "2008ITNS...55.1086D", "2008PhRvD..77h7302H", "2008PhRvD..78k5002K", "2008PhRvL.101w1301F", "2009APh....31..270A", "2009ITNS...56.2989S", "2009JHEP...07..050M", "2009JHEP...07..051R", "2009PhLB..671..391P", "2009PhRvC..79d5807A", "2009PhRvD..79a5014A", "2009PhRvD..80a5003E", "2009PhRvD..80g5018B", "2009PhRvD..80i5024B", "2009PhRvD..80k5019Y", "2010ApPhL..96h3505S", "2010EPJC...67...39B", "2010PhRvD..81e4025F", "2010PhRvD..82c4005B", "2010PhRvD..82e6001C", "2010arXiv1004.0691E", "2011JAP...109h4507G", "2011JAP...110h1301S", "2011PhRvL.107e1301A", "2012ApPhL.101e2601K", "2012EPJC...72.2061S", "2012ExA....34...43F", "2012JCAP...05..034C", "2012JHEP...08..029M", "2012MPLA...2730016E", "2012OExpr..20.1503M", "2012PDU.....1...32G", "2012PhRvD..85f3503L", "2012PhRvD..85g6007E", "2012PhRvL.109b1301E", "2013JCAP...08..041B", "2013JHEP...11..167E", "2013JHEP...11..193E", "2013JInst...8R4001C", "2013NIMPA.716...78A", "2013PhRvD..87k5001G", "2013PhRvD..88k4015I", "2013PhRvL.110x9901A", "2013arXiv1310.8327C", "2013arXiv1311.0029E", "2014IJMPA..2930070K", "2014PhRvD..89g5008G", "2014PhRvD..89h3508N", "2014PhRvL.113q1301H", "2014PhRvL.113q1802B", "2014arXiv1406.3028B", "2015NIMPA.787...68B", "2015PSSBR.252..804M", "2015PhRvD..91e5006K", "2015PhRvD..91i4026I", "2015PhRvD..92h3517L", "2015PhRvL.115b1301H", "2015PhRvL.115y1301I", "2015arXiv150902910T", "2016APh....84...70A", "2016EPJC...76...25A", "2016JCAP...04..027A", "2016JHEP...05..046E", "2016JHEP...06..011Y", "2016JHEP...08..057H", "2016JInst..11P0003L", "2016JPhCS.718d2021F", "2016NIMPA.824...74A", "2016PhRvD..94g3009K", "2016PhRvD..94i2001A", "2016PhRvL.116a1301H", "2016PhRvL.116g1301A", "2016PhRvL.116p1302A", "2016PhRvL.117l1302S", "2017PhLB..772..239H", "2017PhRvD..95c2006B", "2017PhRvD..95h2002A", "2017PhRvL.118p1801A" ]
[ "10.1103/PhysRevD.96.016026", "10.48550/arXiv.1607.01009" ]
1607
1607.01009_arXiv.txt
\vskip -3mm Dark matter (DM) with a mass in the MeV--GeV range is phenomenologically viable and has received increasing attention in recent years~\cite{Essig:2011nj,Essig:2015cda, Graham:2012su, Lee:2015qva, Essig:2012yx, Hochberg:2015pha,*Hochberg:2015fth,*Schutz:2016tid,*Hochberg:2016ntt, Essig:2013lka,*Bird:2004ts,*Borodatchenkova:2005ct,*McElrath:2005bp, *Fayet:2006sp, *Bird:2006jd, *Kahn:2007ru,*Fayet:2007ua,*Essig:2009nc,*Bjorken:2009mm,*Reece:2009un,*Fayet:2009tv,*Yeghiyan:2009xc,*Badin:2010uh,*Echenard:2012iq,*MarchRussell:2012hi,*Essig:2013vha, *Essig:2013goa,*Boehm:2013jpa,*Nollett:2013pwa,*Andreas:2013lya,*Izaguirre:2013uxa,*Battaglieri:2014qoa, *Izaguirre:2014bca,*Batell:2014mga,*Kahn:2014sra,*Krnjaic:2015mbs,*Batell:2009di,*Izaguirre:2015yja, Boehm:2003hm,*Strassler:2006im,*ArkaniHamed:2008qn,*Pospelov:2008jd,*Hooper:2008im,*Feng:2008ya,*Morrissey:2009ur,*Essig:2010ye,*Cohen:2010kn,*Lin:2011gj,*Chu:2011be,*Hochberg:2014dra,*Hochberg:2014kqa}. An important probe for DM is with {\it direct detection} experiments, in which a DM particle in the Milky-Way halo interacts with some target material in a detector, producing an observable signal in the form of heat, phonons, electrons, or photons~\cite{Cushman:2013zza}. The traditional technique of searching for nuclear recoils loses sensitivity rapidly for DM masses below a few GeV, since the DM is unable to transfer enough of its energy to the nucleus, resulting in no observable signal above detector thresholds. However, DM scattering off electrons, whose mass is much less than a nucleus, can lead to observable signals for masses well below 1~GeV~\cite{Essig:2011nj}, opening up vast new regions of parameter space for experimental exploration. DM-electron scattering in direct detection experiments has been investigated for noble liquid targets~\cite{Essig:2011nj,Essig:2012yx} and was demonstrated explicitly to have sensitivity down to DM masses of a few MeV and cross-sections of $\sim 10^{-37}~\rm{cm}^2$~\cite{Essig:2012yx} using published XENON10 data~\cite{Angle:2011th}. Semiconductor targets like silicon (Si) and germanium (Ge) could probe potentially several orders of magnitude of unexplored DM parameter space for masses as low as a few hundred keV~\cite{Essig:2011nj,Graham:2012su,Lee:2015qva,Essig:2015cda}. The feasibility of the required detector technology to detect small ionization signals still needs to be demonstrated and may become available in the next few years, e.g.~with SuperCDMS~\cite{Agnese:2015nto} and DAMIC~\cite{Moroni:2011xs}. In the future, even lower masses could be probed using superconductors or superfluids~\cite{Hochberg:2015pha,Hochberg:2015fth,Schutz:2016tid}. In this letter, we explore using a \emph{scintillator} as the target material to search for dark matter with masses as low as a few hundred keV. One or more scintillation photons are emitted when an electron excited by a DM-electron scattering interaction relaxes to the ground state~\cite{Essig:2011nj}\footnote{Note that~\cite{Starkman:1994gf} proposed the search of one or more photons from Weak-scale dark matter through atomic excitations.}. Scintillation photons with an energy of $\mathcal{O}$(few~eV) could be detected by an array of transition edge sensors (TES) or microwave kinetic inductance detectors (MKIDs) operated at cryogenic temperatures, which surround a scintillating target of volume $\sim \mathcal{O}(({\rm few~ cm})^3)$. The development of such a large array of photodetectors sensitive to single photons is an active area of research~\cite{Pyle}. The target itself should be cooled to cryogenic temperatures to avoid excitations induced by thermal fluctuations and large thermal gradients between it and the detector array. Several good scintillating materials exist. In this letter, we focus on three crystals, sodium iodide (NaI), cesium iodide (CsI), and gallium arsenide (GaAs). Other materials will be mentioned briefly. \vspace{-4mm}
16
7
1607.01009
We suggest a novel experimental concept for detecting MeV-to-GeV-mass dark matter, in which the dark matter scatters off electrons in a scintillating target and produces a signal of one or a few photons. New large-area photodetectors are needed to measure the photon signal with negligible dark counts, which could be constructed from transition edge sensor (TES) or microwave kinetic inductance detector (MKID) technology. Alternatively, detecting two photons in coincidence may allow the use of conventional photodetectors like photomultiplier tubes. We describe why scintillators may have distinct advantages over other experiments searching for a low ionization signal from sub-GeV dark matter, as there are fewer potential sources of spurious backgrounds. We discuss various target choices, but focus on calculating the expected dark matter-electron scattering rates in three scintillating crystals: sodium iodide (NaI), cesium iodide (CsI), and gallium arsenide (GaAs). Among these, GaAs has the lowest band gap (1.52 eV) compared to NaI (5.9 eV) or CsI (6.4 eV), which in principle allows it to probe dark matter masses as low as ∼0.5 MeV , compared to ∼1.5 MeV with NaI or CsI. We compare these scattering rates with those expected in silicon (Si) and germanium (Ge). The proposed experimental concept presents an important complementary path to existing efforts, and its potential advantages may make it the most sensitive direct-detection probe of dark matter down to MeV masses.
false
[ "dark matter", "dark matter masses", "sub-GeV dark matter", "negligible dark counts", "fewer potential sources", "MeV masses", "spurious backgrounds", "NaI", "kinetic inductance detector", "cesium iodide", "sodium iodide", "the expected dark matter-electron scattering rates", "the dark matter", "transition edge sensor", "photomultiplier tubes", "GaAs", "Ge", "CsI", "conventional photodetectors", "MeV" ]
7.814362
-2.235919
54
12501713
[ "Chakravarty, Girish Kumar", "Dey, Ujjal Kumar", "Lambiase, Gaetano", "Mohanty, Subhendra" ]
2016PhLB..763..501C
[ "Plateau inflation in R-parity violating MSSM" ]
8
[ "Physical Research Laboratory, Ahmedabad-380009, India", "Physical Research Laboratory, Ahmedabad-380009, India", "Dipartimento di Fisica, \"E.R. Caianiello\" Universitá di Salerno, I-84084 Fisciano (Sa), Italy", "Physical Research Laboratory, Ahmedabad-380009, India" ]
[ "2016PhRvD..94f3518D", "2017arXiv170703853C", "2018EPJC...78..341K", "2018IJMPA..3350127G", "2019arXiv190907217A", "2020arXiv201215256B", "2021EPJC...81..804P", "2022GrCo...28....1A" ]
[ "astronomy", "physics" ]
2
[ "Inflation", "Supersymmetry", "Supergravity", "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
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[ "10.1016/j.physletb.2016.10.081", "10.48550/arXiv.1607.06904" ]
1607
1607.06904_arXiv.txt
There are atleast two experimental sectors which hint at the existence of scalar fields beyond the Higgs field of the Standard Model (SM). In order to explain the observed anisotropy of the cosmic microwave background (CMB) temperature at the super horizon scales \cite{Ade:2015xua} and at the same time the low value of the tensor-to-scalar ratio $r <0.07$ \cite{Ade:2015xua, Array:2015xqh} one requires an inflaton field with a plateau potential \cite{Starobinsky:1979ty, Starobinsky:1980te, Ellis:2013xoa, Ellis:2014gxa, Kallosh:2013yoa, Kallosh:2015lwa, Chakravarty:2016fin}. The other hint for a scalar field is the 750 GeV diphoton excess which may have been observed by ATLAS \cite{Aaboud:2016tru} and CMS \cite{Khachatryan:2016hje} collaborations which has launched a large number of models which explain the 750 GeV diphoton resonance in the context of left-right models \cite{Dey:2015bur, Dasgupta:2015pbr, Dev:2015vjd, Deppisch:2016scs, Borah:2016uoi, Hati:2016thk, Agarwalla:2016rmw, Ghosh:2016lnu}, Grand Unification \cite{Patel:2015ulo, Aydemir:2016qqj, Nilles:2016bjl} and supersymmetry (SUSY) \cite{Chakraborty:2015gyj, Ding:2015rxx, Allanach:2015ixl, Choudhury:2016jbc, Djouadi:2016oey, Dreiner:2016wwk} and other exotic models (reviewed in \cite{Staub:2016dxq, Strumia:2016wys}). Cosmological implications of the 750 GeV scalar have been studied in \cite{Dhuria:2015ufo, Marzola:2015xbh, Ge:2016xcq, McDonald:2016cdh, Dimopoulos:2016tzn}. Sadly, the signal no longer persists as shown by the updated analysis of $\sqrt{s} = 13$ TeV data of ATLAS and CMS~\cite{Lenzi:2207286, CMS-PAS-EXO-16-027}. In this paper we construct a plateau inflation model in the context of the minimal supersymmetric standard model which can be tested at the LHC in particular where the inflaton can be observed as a TeV scale diphoton resonance. We find that the most economical model which can explain both the phenomenon is to identify the left-handed sneutrino in a $R$-parity violating MSSM as the inflaton and as the diphoton resonance. The identification of the tau-sneutrino as the diphoton resonance has been made in the $R$-parity violating MSSM in \cite{Ding:2015rxx, Allanach:2015ixl}. On the other hand inflation with the singlet right-handed sneutrino has been well studied~\cite{Murayama:1992ua, Murayama:1993xu, Hamaguchi:2001gw, Ellis:2003sq, Antusch:2004hd, BasteroGil:2005gy, Kadota:2005mt} and in MSSM the Higgs-sneutrino inflation along flat-directions \cite{Allahverdi:2005mz, Allahverdi:2006iq, Antusch:2010va, Antusch:2010mv, Pallis:2011ps, Haba:2011uz, Kim:2011ay, Khalil:2011kd, Aulakh:2012st, Arai:2012em, Nakayama:2013nya, Evans:2015mta, Deen:2016zfr} has also been studied. In this paper we consider a supergravity model with no-scale like K\"ahler potential and a superpotential which includes $R$-parity violating non-renormalizable operators at all orders. In this model we consider the supersymmetry breaking occurs via a Polonyi field which takes a non-zero vacuum expectation value after the end of inflation in the present epoch. The proper choice of superpotential in Polonyi field leads to much heavier Polonyi mass compared to gravitino mass to avoid the cosmological moduli problem and to obtain the vanishingly small cosmological constant $~10^{-120}$ \cite{Coughlan:1983ci, Ellis:1986zt, Goncharov:1984qm, Linde:2011ja, Dudas:2012wi}. The SUSY breaking generates masses which are of the TeV scale for all the SUSY scalar partners (like squarks, sneutinos and sleptons). The TeV scale sleptons are used in the loops for the production and decay of the TeV scale sneutrino. The production and decay vertices which involve sneutrino-quark and sneutrino-sleptons are generated by the SUSY breaking by the hidden sector Polonyi field. The rest of the paper is organized as follows. In Sec.~\ref{sec:model} we introduce the relevant K\"ahler potential and superpotential and construct the $D$-term and $F$-term potentials. In Sec.~\ref{sec:inflationDflat} we choose the $D$-flat Higgs-sneutrino direction that gives the required plateau potential from the $F$-term. We show how interaction terms arise from the Polonyi field SUSY breaking which ultimately gives rise to the diphoton production and decay vertices in Sec.~\ref{sec:soft}. We then apply the model to the calculation of the $\sim$ TeV sneutrino production and decay and show the cross section for the diphoton resonance which can be a tentative signature for this model in Sec.~\ref{sec:diphoton}. We conclude and list future implications of the model in Sec.~\ref{sec:concl}.
\label{sec:concl} Plateau inflation in MSSM with a $D$-flat combination of Higgs fields has been studied in \cite{Chakravarty:2016fin}. In this paper we show that plateau inflation can be achieved in the $R$-parity violating MSSM where the TeV scale sneutrino and charged slepton masses can give testable LHC prediction like a diphoton resonance with significant cross section ($\sigma_{pp \rightarrow \phi \rightarrow \gamma \gamma} \sim 10 ~{\rm fb}$). This signal may show up as TeV scale resonance in the future. Sneutrino inflation models also have applications in leptogenesis by the Affleck-Dine mechanism \cite{Affleck:1984fy} as has been studied earlier \cite{Garcia:2013bha} and these leptogenesis mechanisms can be studied in our specific model of sneutrino plateau inflation.
16
7
1607.06904
Inflation with plateau potentials give the best fit to the CMB observables as they predict tensor to scalar ratio stringently bounded by the observations from Planck and BICEP2/Keck. In supergravity models it is possible to obtain plateau potentials for scalar fields in the Einstein frame which can serve as the inflation potential by considering higher dimensional Planck suppressed operators and by the choice of non-canonical Kähler potentials. We construct a plateau inflation model in MSSM where the inflation occurs along a sneutrino-Higgs flat direction. A hidden sector Polonyi field is used for the breaking of supersymmetry after the end of the inflation. The proper choice of superpotential leads to strong stabilization of the Polonyi field, m<SUB>Z</SUB><SUP>2</SUP> ≫m<SUB>3/2</SUB><SUP>2</SUP>, which is required to solve the cosmological moduli problem. Also, the SUSY breaking results in a TeV scale gravitino mass and scalar masses and gives rise to bilinear and trilinear couplings of scalars which can be tested at the LHC. The sneutrino inflation field can be observed at the LHC as a TeV scale diphoton resonance like the one reported by CMS and ATLAS.
false
[ "non-canonical Kähler potentials", "scalar fields", "plateau potentials", "scalar ratio", "higher dimensional Planck", "Inflation", "Polonyi field", "ATLAS", "Planck", "the inflation potential", "TeV", "a TeV scale diphoton resonance", "CMS", "strong stabilization", "The sneutrino inflation field", "Kähler", "a TeV scale gravitino mass and scalar masses", "operators", "a plateau inflation model", "tensor" ]
10.528646
-1.285353
89
12462322
[ "Dalmasse, Kevin", "Nychka, Douglas", "Gibson, Sarah", "Fan, Yuhong", "Flyer, Natasha" ]
2016FrASS...3...24D
[ "ROAM: a Radial-basis-function Optimization Approximation Method for diagnosing the three-dimensional coronal magnetic field" ]
9
[ "National Center for Atmospheric Research, , Boulder, USA", "-", "-", "-", "-" ]
[ "2017FrASS...4....3G", "2017SSRv..212.1345C", "2018ApJ...852..141D", "2019ApJ...877..111D", "2019ApJ...883...55Z", "2019sfsw.book...41C", "2020A&A...642A...2R", "2020ApJ...898...70M", "2021ApJ...907...23C" ]
[ "astronomy" ]
3
[ "solar corona", "Solar magnetic field", "infrared", "statistical methods", "Radial basis functions", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1908ApJ....28..315H", "1924ZPhy...30...93H", "1977A&A....59..223S", "1981A&A...100..197A", "1982SoPh...78..157B", "1985SoPh...98..341B", "1989A&A...221..326D", "1997SoPh..174...31W", "1998JGR...10314511G", "1999ApJ...522..524C", "2000ApJ...540.1150W", "2000SoPh..195...89Y", "2001ASPC..248..597F", "2003rbf..book.....B", "2004ApJ...612..519V", "2004ApJ...613L.177L", "2004SoPh..219...87W", "2005A&A...433..335V", "2006A&A...446..691A", "2006ApJ...651.1229J", "2007ApJ...662..677J", "2007SoPh..245..263V", "2008SoPh..247..411T", "2009ASPC..405..429R", "2011ApJ...727..101J", "2012ASPC..463..227R", "2012ApJ...747...65I", "2012ApJ...756..153M", "2012ApJ...758...60F", "2013SoPh..288..617R", "2014ApJ...790..163T", "2014ApJ...792...23P", "2014SoPh..289.2927K", "2015ASSL..415.....V", "2016FrASS...3....8G", "2016JCoPh.316...39F", "2016JGRA..121.7470T", "2017SSRv..210....5S" ]
[ "10.3389/fspas.2016.00024", "10.48550/arXiv.1607.03460" ]
1607
1607.03460_arXiv.txt
\label{sec:S-Introduction} Modification to the polarization of light is one of the many signatures of a non-zero magnetic field in the solar corona, and more generally, in the solar atmosphere \citep[\eg][and references therein]{Stenflo15}. Several mechanisms producing or modifying the polarization of light have been observed and studied in the solar corona at different wavelengths including, but not limited to, the Zeeman and Hanle effects \citep[see \eg][and references therein]{Hale08,Hanle24,Bird85,White97,Casini99,Lin04,Gibson16}. The former induces a frequency-modulated polarization while the latter induces a depolarization of scattered light \citep[\eg][]{SahalBrechot77,Bommier82,Rachmeler13,LopezAriste15}. Both mechanisms allow us to probe the strength and direction of the coronal magnetic field. Coronal polarization associated with these two mechanisms is currently measured above the solar limb by the Coronal Multichannel Polarimeter from forbidden coronal lines such as the Fe XIII lines \citep[10747 $\AA$ and 10798 $\AA$;][]{Tomczyk08}. For these two lines, the circular polarization signal is dominated by the Zeeman effect while the linear polarization signal is dominated by the Hanle effect \citep[\eg][]{Judge06}. Translating the polarization maps of CoMP into magnetic field maps is a challenging task. The main difficulty is that the solar corona is optically thin at these wavelengths \citep[\eg][]{Rachmeler12,Plowman14}. The observed signal is therefore the integrated emission of all the plasma along the {line of sight (LOS).} Hence, the polarization maps cannot, in general, be directly inverted into 2D maps of the plane-of-sky (POS) magnetic field. On the other hand, extracting individual magnetic information at specific positions along the LOS is extremely difficult without stereoscopic observations \citep[\eg][]{Kramar14}. Another limitation is that the Hanle effect associated with the aforementioned forbidden infrared lines operates in the saturated regime \citep[\eg][]{Casini99,Tomczyk08}. Accordingly the linear polarization signal measured by CoMP is sensitive to the direction of the magnetic field but not its strength. Deriving the magnetic field associated with the polarization maps of CoMP therefore requires a different approach than the single point inversion that can be done with, \eg photospheric polarimetric measurements. The alternate approach we propose to follow is to combine a parameterized 3D magnetic field model with forward modeling of the polarization signal observed by CoMP. For that purpose, we take advantage of the Coronal Line Emission (CLE) polarimetry code developed by \cite{Casini99} and integrated into the FORWARD package. FORWARD\footnote{\url{http://www.hao.ucar.edu/FORWARD/}} is a Solar Soft\footnote{\url{http://www.lmsal.com/solarsoft/}} IDL package designed to perform forward modeling of various observables including, \eg visible/IR/UV polarimetry, EUV/X-ray/radio imaging, and white-light coronagraphic observations \citep{Gibson16}. The goal is then to optimize a user-specified likelihood {function} comparing the polarization signal predicted by FORWARD to the real one and find the parameters of the magnetic field model such that the predicted signal fits the real data. In the present paper, we develop and test a new method for performing fast and efficient optimization in a $d$-dimensional parameter space that may be used for converting the polarization observations of CoMP into magnetic field data. The optimization method, called ROAM (Radial-basis-functions Optimization Approximation Method) is designed to be general enough so that it can be applied independently of the dimension and size of the parameter space, the 3D magnetic field model, the type of observables (provided that one can forward model them), and the form of the likelihood {function} used for comparing the predicted signal to the real one. ROAM is introduced in \sect{S-Method}. \sect{S-Results} describes the results of multiple applications of ROAM to a synthetic test bed as validation of the optimization method. Our conclusions are then summarized in \sect{S-Conclusions}.
\label{sec:S-Conclusions} In this paper, we introduced and validated a new optimization method for model-data fitting, ROAM (Radial-basis-functions Optimization Approximation Method). Our primary motivation for this work has been to develop a novel approach for diagnosing the solar coronal magnetic field by combining a parameterized 3D magnetic field model with forward modeling of coronal polarization. From various tests applied to the synthetic test bed of a coronal magnetic flux rope, we showed that ROAM allows for fast, efficient, and accurate model-data fitting in a $d$-dimensional parameter space. These test cases further enabled us to analyze and specify a framework for an optimal application of ROAM. Applying our method with forward modeling of IR coronal polarimetry, we demonstrated that ROAM can be exploited for converting coronal polarimetric measurements into magnetic field data. The use of our model-data fitting method therefore opens new perspectives for the development and exploitation of coronal polarimetric measurements such as the ones routinely performed by CoMP \citep{Tomczyk08} and future telescopes such as the Daniel K. Inoue Solar Telescope\footnote{\url{http://www.ifa.hawaii.edu/~schad/dlnirsp/}} and the Coronal Solar Magnetism Observatory (Tomczyk et al., submitted), but also for a wider range of coronal observations including, \eg UV \citep[see \eg][]{Fineschi01,Raouafi09} and radio polarimetry (\eg \citeauthor{White97}, \citeyear{White97}; \citeauthor{Gelfreikh04}, \citeyear{Gelfreikh04}; see also \citeauthor{Gibson16}, \citeyear{Gibson16}, for discussion of multiwavelength magnetometry). Beyond the analysis of coronal polarimetric measurements, ROAM offers interesting perspectives for magnetic field reconstruction models. Most of the current 3D diagnostics of the coronal magnetic field of solar active regions (ARs) are derived from the analysis of magnetic field reconstruction models including, \eg force-free field extrapolations of the photospheric magnetic field \citep[see \eg][and references therein]{Alissandrakis81,Demoulin89,Wheatland00,Yan00,Wiegelmann04,Amari06,Malanushenko12}, {and} magneto-frictional methods \citep[see \eg][and references therein]{VanBallegooijen04,Valori05,Valori07,Jiang11,Inoue12,Titov14}. ROAM could, in principle, be used to perform model-data fitting with such reconstruction models that either already are (\ie through the poloidal and axial flux for the {magneto-frictional methods with} flux rope insertion) or could be (\eg through the photospheric force-free parameter for both force-free field extrapolations and magneto-frictional methods {without flux rope insertion)} parameterized. {The extensive work performed} over the years in terms of forward modeling of various observables \citep[see \eg][and references therein]{Gibson16} would then allow for using several types of different observations to constrain the parameters of the magnetic field reconstruction models. ROAM therefore opens new perspectives for including coronal polarimetric measurements into magnetic field reconstructions and, more generally, for data-optimized reconstruction of the solar coronal magnetic field. Such perspectives will be tackled in the framework of the Data Optimized Coronal Field Model\footnote{\url{http://www.hao.ucar.edu/DOCFM/}} (DOCFM), a collaborative project that will make use of ROAM. Finally, we wish to mention that ROAM is not limited to coronal magnetic field diagnostics and could be used for other optimization problems. The method will be of particular interest for model-data fitting for which a model evaluation (here, the evaluation of the model itself and/or the forward modeling of an observable if applicable) is computationally expensive.
16
7
1607.03460
The Coronal Multichannel Polarimeter (CoMP) routinely performs coronal polarimetric measurements using the Fe XIII 10747 Å and 10798 Å lines, which are sensitive to the coronal magnetic field. However, inverting such polarimetric measurements into magnetic field data is a difficult task because the corona is optically thin at these wavelengths and the observed signal is therefore the integrated emission of all the plasma along the line of sight. To overcome this difficulty, we take on a new approach that combines a parameterized 3D magnetic field model with forward modeling of the polarization signal. For that purpose, we develop a new, fast and efficient, optimization method for model-data fitting: the Radial-basis-functions Optimization Approximation Method (ROAM). Model-data fitting is achieved by optimizing a user-specified log-likelihood function that quantifies the differences between the observed polarization signal and its synthetic/predicted analogue. Speed and efficiency are obtained by combining sparse evaluation of the magnetic model with radial-basis-function (RBF) decomposition of the log-likelihood function. The RBF decomposition provides an analytical expression for the log-likelihood function that is used to inexpensively estimate the set of parameter values optimizing it. We test and validate ROAM on a synthetic test bed of a coronal magnetic flux rope and show that it performs well with a significantly sparse sample of the parameter space. We conclude that our optimization method is well-suited for fast and efficient model-data fitting and can be exploited for converting coronal polarimetric measurements, such as the ones provided by CoMP, into coronal magnetic field data.
false
[ "coronal magnetic field data", "magnetic field data", "coronal polarimetric measurements", "such polarimetric measurements", "parameter values", "the coronal magnetic field", "a parameterized 3D magnetic field model", "a coronal magnetic flux rope", "Optimization Approximation Method", "ROAM", "sight", "sparse evaluation", "the observed polarization signal", "forward modeling", "Model-data fitting", "the magnetic model", "the polarization signal", "the log-likelihood function", "a user-specified log-likelihood function", "Radial" ]
12.401613
15.609333
2
12473496
[ "Ricotti, Massimo" ]
2016MNRAS.462..601R
[ "X-ray twinkles and Population III stars" ]
33
[ "Department of Astronomy, University of Maryland, College Park, MD 20742, USA" ]
[ "2016PASA...33...51L", "2017A&A...602A.103W", "2017ApJ...838..117N", "2017MNRAS.464.1365M", "2018MNRAS.473.5308M", "2018MNRAS.474..443G", "2018MNRAS.478.5591M", "2018MNRAS.479..667R", "2018MNRAS.479.4544M", "2019MNRAS.483.1582H", "2019MNRAS.489.1880H", "2020MNRAS.492.4858H", "2020MNRAS.495.2966S", "2021MNRAS.508.1954S", "2021MNRAS.508.2784W", "2021MNRAS.508.6176P", "2021MNRAS.508.6193P", "2023ARA&A..61...65K", "2023ApJ...959...17S", "2023MNRAS.521.4039S", "2023MNRAS.521.5334P", "2023MNRAS.522.2495G", "2023MNRAS.524.2290N", "2023MNRAS.525..428H", "2024ApJ...962...49B", "2024ApJ...962...62F", "2024ApJ...964...62Z", "2024BSRSL..93..700A", "2024MNRAS.527.5023L", "2024MNRAS.528.6895P", "2024PhRvD.109j3026Q", "2024arXiv240304824S", "2024arXiv240609483Q" ]
[ "astronomy" ]
2
[ "stars: Population III", "supernovae: general", "early Universe", "X-rays: general", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1985ApJ...298..268S", "1997ApJ...474....1T", "1998A&A...335..403G", "1999ApJ...527L...5B", "2000ApJ...534...11H", "2001ApJ...553..499O", "2001ApJ...560..580R", "2001ApJ...563....1V", "2001MNRAS.328..969B", "2002ApJ...567..532H", "2002ApJ...575...33R", "2002ApJ...575...49R", "2002ApJ...576..653S", "2002Sci...295...93A", "2003MNRAS.338..273M", "2003Natur.422..871U", "2004ApJ...611...40C", "2004MNRAS.352..547R", "2005ApJ...629..259R", "2005MNRAS.357..207R", "2006MNRAS.369.1719M", "2007ApJ...654..897B", "2008ApJ...680..829R", "2008ApJ...685...21R", "2008ApJ...685...40W", "2008MNRAS.387L...8V", "2009ApJ...693.1859B", "2009MNRAS.392...45R", "2009MNRAS.392L..45R", "2009Sci...325..601T", "2010ApJ...714..287G", "2010MNRAS.403...45S", "2011ApJ...732..114L", "2011ApJ...739....2P", "2011ApJ...741...17B", "2011ApJ...741...18B", "2012ApJ...745...68T", "2012ApJ...746..125H", "2012ApJ...747....9P", "2012MNRAS.425.3058C", "2013ApJ...767..163P", "2013ApJ...772..106M", "2013ApJ...776...34K", "2013ApJ...776L..31F", "2014ApJ...780..145T", "2014ApJ...791..110X", "2014MNRAS.440.3778J", "2014MNRAS.442.2560W", "2014MNRAS.445..581H", "2015ApJ...802L..19R", "2015ApJ...805..130K", "2015ApJ...813..109D", "2015MNRAS.452.1152J", "2015MNRAS.453.1305W", "2016A&A...594A..13P", "2016MNRAS.455..282H", "2016MNRAS.462.1164H", "2016PhRvL.116f1102A" ]
[ "10.1093/mnras/stw1672", "10.48550/arXiv.1607.04289" ]
1607
1607.04289_arXiv.txt
The number of first stars (\popIII) per comoving volume that forms in the early universe determines the level and homogeneity of metal pre-enrichment of the intergalactic medium (IGM). Metal pre-enrichment is important for modeling the formation of the first dwarf galaxies and predict the number of pre-reionization fossils in the Local Group \citep{RicottiG:05, BovillR:09}. There are two main approaches widely used for modeling the formation of the first dwarf galaxies in cosmological simulations: (a) metal enrichment is calculated self-consistently resolving the formation of Pop~III stars at $z>10$ in relatively small (1-4 cMpc$^3$) cosmological volumes \citep{RicottiGSa:02, RicottiGS:08,Wise:08, Wise:14, Muratov:13a}; (b) a metallicity floor (typically $Z \sim 10^{-3}$~Z$_\odot$) is introduced everywhere in the IGM in order to run zoom simulations of dwarf galaxies from high-redshift to the present \citep{Gnedin:10, Tassis:12, Christensen:12, Kuhlen:13, Hopkins:14, Thompson:14, Wheeler:15}. The second method is not suited for capturing global feedback loops that might affect the local metallicity floor and the intensity of the radiation backgrounds, which are both important in determining the fraction of dark matter halos that remain dark. However, the self-consistent method in (a) also have limitations: \begin{enumerate} \item The gravitational potential of dark matter halos drives the collapse of proto-\popIII stars until the gas becomes self-gravitating at scales of a few AU \citep{Bromm1999, Abel2002}. Hence, in order to capture the formation of Pop~III stars is necessary to resolve the gravitational potential at the center of the minihalos of mass $10^5$~M$_\odot$ with at least several tens of particles. A dark matter resolution of about $100$~M$_\odot$ is required, setting a limit on the cosmological volume that can be simulated (e.g., $512^3$ simulation with $m_p=100$ M$_\odot$ has volume of 3~Mpc$^3$). \item On the other hand, the small cosmological volume required to achieve the resolution necessary to resolve \popIII star formation in the smallest dark matter halos, prevents a self-consistent calculation of the radiation backgrounds that are important in determining the number of \popIII stars, especially in underdense regions where local feedback effects are sub dominant. \end{enumerate} One can choose to calculate the self-consistent backgrounds even though the simulated volume is too small for numerical convergence, or include a tabulated external background from analytical models. In the first case, as soon as the very first Pop~III star is created, the dissociating radiation background jumps from zero to a sufficiently large value to destroy very rapidly all relic \H2 in the IGM \citep[\eg,][]{RicottiGSb:02}. For the second choice, often a tabulated background \citep[\eg,][]{HaardtMadau:12} is adopted. But in the case of the formation of the Pop~III stars at $z=40-10$, the use of backgrounds derived from observations at $z<10$ is not justified. Hence, both choices are not satisfactory. In this paper we use simple analytical calculations to estimate self-consistently the number of \popIII stars in the early universe and the radiation background they produce during their short life on the main sequence and by their SN remnants. We also consider other sources of X-rays: accreting intermediated mass black holes (IMBHs), high mass X-ray binaries (HMXRBs) \citep{Xu:14, Jeon:14, Jeon:15}, and miniquasars. We derive what is the level of X-ray emissivity that maximizes the number of \popIII stars forming at $z \simgt 10-15$. The basic idea is simple as noted by several authors before \citep{Oh2001,Venkatesan2001, Machacek2003,RicottiO:04}. An X-ray background can both suppress \popIII star formation in the smallest minihalos due to IGM heating (increasing the Jeans mass in the IGM) and promote \popIII star formation by increasing the gas electron fraction in gas collapsed into minihalos and thus promoting H$_2$ formation via the catalyst $H^-$. The number of \popIII stars depends on the minimum dark matter halo mass, $M_{\rm cr}$, in which a \popIII star can form as a function of redshift; the smaller the critical mass the more numerous the \popIII stars. However, $M_{\rm cr}$ depends on the X-ray background, as explained above, and the H$_2$ dissociating background (UV in the Lyman-Werner bands); since \popIII stars are responsible for producing the dissociating and X-ray backgrounds, a feedback loop is in play. The model is presented in \S~\ref{sec:model} and the results in \S~\ref{sec:res}. The discussion is in \S~\ref{sec:disc}, and summary and conclusions are in \S~\ref{sec:sum}. We use Planck cosmology $(\Omega_m, \Omega_\lambda, \Omega_b, h, n_s, \sigma_8)=$(0.308, 0.692, 0.0482, 0.678, 0.968, 0.829) \citep{Planck2015}.
\label{sec:sum} A low level of X-ray emission in the early universe, although has a negligible effect on reionization and the optical depth to Thompson scattering, goes a long way in enhancing the number of \popIII stars and dwarf galaxies with halo masses $M_{halo}<10^8$~M$_\odot$ that can only form before IGM reheating and reionization at redshift $z \sim 6-10$. The maximum number of \popIII stars is obtained when the critical halo mass in which gas can cool in less than a Hubble time equals the Jeans mass of the IGM. This happens for $K_X=0.01 K_{LW}^{0.5}$, where $K_X$ and $K_{LW}$ are the mean energies of the first sources of light in the soft X-ray and \LW bands in units of $6\times 10^{53}$~ergs, respectively. This low level of X-ray emission does not require assumptions on the presence of unknown X-ray sources such as IMBH or miniquasars: it is necessarily produced by SN explosions and SN remnants of \popIII stars, thus is an unavoidable consequence of \popIII star formation with a top heavy IMF (unless most stars are more massive than $250$~M$_\odot$, that would lead to direct collapse into black holes without SN explosion). In addition, if a non-negligible fraction of \popIII stars explode as hypernovae or PISNe, the soft X-rays from their remnants and explosions is sufficiently large to promote the formation of \popIII stars to about 400 per cMpc$^3$, that is near the maximum number of \popIII stars with typical masses $10-40$~M$_\odot$ that can form in any of our models with different $K_X$. A higher X-ray flux than the one provided by \popIII hypernovae, for instance produced by accretion onto IMBH from \popIII stars and miniquasars, would suppress the number of \popIII stars because of the excessive heating of the IGM. We find that X-rays emitted by HMXRBs have a negligible effect in boosting the number of \popIII stars when compared to the soft X-ray emission from the first SN remnants. The implications of a large number of \popIII stars include: i) a copious production of black holes with masses similar to the ones detected by LIGO ($\sim 10-30$~M$_\odot$) via gravitational wave emission \citep{LIGO2016} (about $10^4$ BHs remnants of \popIII stars are estimated within the Milky Way in the hypernova scenario); ii) would provide supermassive black holes seeds; iii) although \popIII stars cannot fully reionize the IGM due to their bursty nature, they can contribute to the reionization process in a manner similar to an early X-ray background \citep{HartleyRicotti2016}; iv) finally, since the mean distance between minihalos hosting \popIII stars is small ($n_{pop3}^{-1/3} \sim 13$~kpc physical at $z \sim 10$), it is easier for the metals ejected by their SN remnants to fill rather uniformly the IGM. A low-level metal pre-enrichment of the IGM (\ie, the metallicity floor often assumed in zoom simulations of galaxy formation), promotes the formation of pre-reionization dwarf galaxies and increases the number of their fossil relics in the Local Group \citep{RicottiG:05, BovillR:09}. The population of ultra-faint dwarfs discovered since 2005 orbiting the Milky Way \citep{Belokurov2007a, DES2015, Koposov2015} is indeed consistent with a large population of pre-reionization dwarf galaxies \citep{BovillR:11a,BovillR:11b}. \subsection*
16
7
1607.04289
Population III stars are typically massive stars of primordial composition forming at the centres of the first collapsed dark matter structures. Here we estimate the optimal X-ray emission in the early universe for promoting the formation of Population III stars. This is important in determining the number of dwarf galaxies formed before reionization and their fossils in the local universe, as well as the number of intermediate-mass seed black holes. A mean X-ray emission per source above the optimal level reduces the number of Population III stars because of the increased Jeans mass of the intergalactic medium, while a lower emission suppresses the formation rate of H<SUB>2</SUB> preventing or delaying star formation in dark matter minihaloes above the Jeans mass. The build-up of the H<SUB>2</SUB> dissociating background is slower than the X-ray background due to the shielding effect of resonant hydrogen Lyman lines. Hence, the nearly unavoidable X-ray emission from supernova remnants of Population III stars is sufficient to boost their number to few tens per comoving Mpc<SUP>3</SUP> by redshift z ∼ 15. We find that there is a critical X-ray to ultraviolet energy ratio emitted per source that produces a universe where the number of Population III stars is largest: 400 per comoving Mpc<SUP>3</SUP>. This critical ratio is very close to the one provided by 20-40 M<SUB>⊙</SUB> Population III stars exploding as hypernovae. High-mass X-ray binaries in dwarf galaxies are far less effective at increasing the number of Population III stars than normal supernova remnants, we thus conclude that supernovae drove the formation of Population III stars.
false
[ "Population III stars", "star formation", "massive stars", "redshift z", "dark matter minihaloes", "resonant hydrogen Lyman lines", "High-mass X-ray binaries", "normal supernova remnants", "the optimal X-ray emission", "ray", "supernova remnants", "A mean X-ray emission", "intermediate-mass seed black holes", "Jeans", "-", "Mpc", "the first collapsed dark matter structures", "the X-ray background", "dwarf galaxies", "the increased Jeans mass" ]
11.213357
8.58313
164
12467950
[ "Tenkanen, Tommi" ]
2016JHEP...09..049T
[ "Feebly interacting dark matter particle as the inflaton" ]
50
[ "University of Helsinki and Helsinki Institute of Physics, Helsinki, Finland" ]
[ "2017IJMPA..3230023B", "2017JCAP...12..001T", "2017JHEP...11..080C", "2017JHEP...12..072H", "2017PhRvD..95c5004A", "2017PhRvD..96b3001H", "2018EPJC...78...26C", "2018JCAP...02..006E", "2018JCAP...05..036C", "2018JCAP...12..020B", "2018MPLA...3350087L", "2018PhRvD..97f3002H", "2018PhRvD..98j3525M", "2018arXiv180703308H", "2018arXiv180709952P", "2019EPJC...79...99B", "2019EPJC...79..862L", "2019JCAP...03..012B", "2019JCAP...04..035T", "2019JHEP...02..020M", "2019JHEP...05..060C", "2019JHEP...06..112L", "2019NuPhB.94414643Y", "2019PDU....25..317M", "2019PhLB..795..657B", "2019PhRvD..99e5012B", "2019PhRvD.100d3507H", "2019PhRvD.100h3515T", "2019PhRvL.122i1802H", "2019PhRvL.122p1301R", "2020GReGr..52...33T", "2020JCAP...07..030N", "2020PhLB..80935716T", "2020PhRvD.101g5006B", "2020PhRvD.102h3534H", "2020PhRvL.124f1301M", "2021EPJC...81..877B", "2021JCAP...04..037B", "2021PhLB..82136614L", "2021PhRvD.103f3525E", "2021PhRvD.104f3010C", "2021arXiv211214668R", "2022JHEP...09..231G", "2022JHEP...12..105G", "2022PhRvD.106b3506H", "2023JCAP...02..024I", "2023JHEP...09..012H", "2024PhRvD.109a5002Y", "2024PhRvD.109b3526G", "2024PhRvD.109f3037G" ]
[ "astronomy", "physics" ]
2
[ "Cosmology of Theories beyond the SM", "Effective field theories", "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1984qfcs.book.....B", "1999PhR...314....1L", "2002PhRvL..88i1304M", "2008PhLB..659..703B", "2009JCAP...06..029B", "2009PhLB..678....1D", "2009PhRvD..80l3507L", "2010JHEP...03..080H", "2011JCAP...07..025K", "2011JHEP...08..060Y", "2011PhR...497...85M", "2011PhRvD..83l3522L", "2012NJPh...14l5003M", "2013PhRvD..88d3502V", "2013arXiv1311.5297B", "2014JCAP...01..003B", "2014JLTP..176..733M", "2014PDU.....5...75M", "2014arXiv1412.3041B", "2015JCAP...11..001N", "2015JCAP...11..015K", "2015JHEP...03..048E", "2015PhLB..750..194S", "2015PhLB..751..201K", "2015PrPNP..85....1K", "2016A&A...594A..13P", "2016A&A...594A..20P", "2016EPJC...76..289C", "2016JCAP...01..006B", "2016JCAP...03..018B", "2016JCAP...06..022K", "2016JCAP...09..037T", "2016PhRvD..93j3531B", "2016PhRvD..93l3513A", "2016PhRvD..94f3506H", "2016PhRvD..94h3516T" ]
[ "10.1007/JHEP09(2016)049", "10.48550/arXiv.1607.01379" ]
1607
1607.01379_arXiv.txt
Extensions of the Standard Model of particle physics (SM) typically contain many scalar fields. Their role in explaining the observed curvature power spectrum and dark matter (DM) abundance, different early Universe phase transitions, matter-antimatter asymmetry, and many other phenomena have been studied extensively in the literature, as discussed, for example, in the recent reviews \cite{Mazumdar:2010sa,Morrissey:2012db,Martin:2013tda,Klasen:2015uma}. In this work, we study a class of beyond the SM scalar fields to address two major issues in cosmology: inflation and dark matter. During the years 2009--13, the European Space Agency's Planck satellite measured properties of the Cosmic Microwave Background (CMB) and either supported, constrained, or even ruled out many scenarios of the early Universe physics. In particular, the Planck results -- together with many different astrophysical observations at different scales --- have shown overwhelming evidence for the existence of an unknown non-baryonic dark matter component, whose abundance in the Universe is now known to be $\Omega_{\rm DM} h^2\simeq 0.12$ \cite{Ade:2015xua}. How this abundance was produced in the early Universe is however still unknown, as no conclusive dark matter signals have shown up in experiments \cite{Klasen:2015uma}. The Planck satellite placed bounds also on many inflationary scenarios by measuring the spectral index of primordial power spectrum to a high accuracy, $\Delta n_s= 0.0060$, and bounding the tensor-to-scalar ratio to $r<0.11$ \cite{Ade:2015lrj}. Among inflationary models the best fit to the Planck data is provided by different Starobinsky-like models, such as Higgs inflation \cite{Bezrukov:2007ep} or $s$-inflation \cite{Lerner:2009xg,Kahlhoefer:2015jma,Aravind:2015xst}, where a non-minimal coupling between gravity and quantum fields typically plays a crucial role. In this work, we connect a Starobinsky-like inflationary model to dark matter production which occurs at a later stage in the history of the Universe. We consider a scenario where a $Z_2$-symmetric scalar field first drives cosmic inflation, then reheats the Universe but remains out-of-equilibrium itself, and finally comprises the observed dark matter abundance, produced by particle decays \`{a} la freeze-in mechanism \cite{McDonald:2001vt,Hall:2009bx,Yaguna:2011qn,Blennow:2013jba,Dev:2013yza,Elahi:2014fsa, Dev:2014tla,Kang:2015aqa,Nurmi:2015ema,Kainulainen:2016vzv}. As the $Z_2$ symmetric scalar field serves as both the inflaton and a FIMP ('Feebly Interacting Massive Particle') dark matter candidate, we name our scenario as the 'fimplaton' model\footnote{In contrast to our scenario, the standard $s$-inflation model \cite{Lerner:2009xg,Kahlhoefer:2015jma} and several other scenarios \cite{Bastero-Gil:2015lga} could be referred to as 'wimplaton' models, a term coined by the latter reference.}. The paper is organized as follows: in Section \ref{model} we present the model and discuss general aspects of the phenomenology and requirements for the fimplaton scenario. Then, we present the scenario in a chronological order as it may have occured in the history of the Universe: first, in Sections \ref{inflationdynamics} and \ref{inflationobservables}, we study how the fimplaton with a non-minimal coupling to gravity drives inflation, then in Section \ref{reheating} we present a mechanism for reheating the Universe, and in Section \ref{freezein} discuss how this same scalar field comprises the observed DM abundance, produced by decays of other fields. Finally, in Section \ref{conclusions}, we conclude and present an outlook.
\label{conclusions} In this work we have studied a scenario where a $Z_2$-symmetric scalar field, non-minimally coupled to gravity, drives cosmic inflation, reheats the Universe but remains \linebreak out-of-equilibrium itself, and later comprises the observed dark matter abundance, produced by particle decays \`{a} la freeze-in mechanism. As the $Z_2$-symmetric scalar field serves as both the inflaton and a FIMP ('Feebly Interacting Massive Particle') dark matter candidate, we have named our scenario as the 'fimplaton' model. Because we wanted to work as model-independently as possible, we did not specify the fimplaton's connection to the known Standard Model physics nor its interactions besides its self-interaction coupling, $\lambda_{\rm \phi}\phi^4$, non-minimal coupling to gravity, $\xi_{\rm \phi}\phi^2R$, and coupling to another scalar field, $g\phi^2\sigma^2$. It would be interesting to study whether already e.g. a simple Higgs portal model, where $\sigma$ is the SM Higgs % and $\phi$ a portal scalar, could accommodate the fimplaton scenario. We have shown the fimplaton model constitutes an interesting example of a scenario where even very small couplings can be responsible for both inflation, reheating, and production of the observed dark matter abundance. Although the (somewhat conservative) coupling and mass windows where the model works are relatively narrow,\linebreak $10^{-9}\lesssim \lambda_{\rm \phi}\lesssim g\lesssim 10^{-7}$, $3 \rm{keV} \lesssim m_{\rm \phi}\lesssim 85 \rm{MeV}$, the scenario is shown to provide a successful connection between cosmic inflation and dark matter abundance. Furthermore, as shown in Section \ref{inflationobservables}, the model may be distinguishable from other inflationary models of the same type, namely the Higgs inflation and $s$-inflation, by the next generation CMB satellites. An interesting aspect of the fimplaton model is that to produce the observed curvature perturbation amplitude within the scenario, the non-minimal coupling has to take a relatively small value, $\xi_{\rm \phi}=\mathcal{O}(1)$. This is indeed a very small value, as Higgs and $s$-inflation models typically require $\xi=\mathcal{O}(10^4)$. As quantum corrections in a curved background have been shown to generate small non-minimal couplings even if they are initially set to zero, it would be particularly interesting to apply the fimplaton scenario to concrete model setups. As SM extensions typically contain many new scalar fields, studies of their role in both inflation and dark matter production together provide many new ways to extract information about SM extensions and physics of the early Universe in general.
16
7
1607.01379
We present a scenario where a Z <SUB>2</SUB>-symmetric scalar field ϕ first drives cosmic inflation, then reheats the Universe but remains out-of-equilibrium itself, and finally comprises the observed dark matter abundance, produced by particle decays à la freeze-in mechanism. We work model-independently without specifying the interactions of the scalar field besides its self-interaction coupling, λϕ <SUP>4</SUP>, non-minimal coupling to gravity, ξϕ <SUP>2</SUP> R, and coupling to another scalar field, gϕ <SUP>2</SUP> σ <SUP>2</SUP>. We find the scalar field ϕ serves both as the inflaton and a dark matter candidate if 10<SUP>-9</SUP> ≲ λ ≲ g ≲ 10<SUP>-7</SUP> and 3 keV ≲ m<SUB> ϕ </SUB> ≲ 85 MeV for ξ ={O}(1) . Such a small value of the non-minimal coupling is also found to be of the right magnitude to produce the observed curvature perturbation amplitude within the scenario. We also discuss how the model may be distinguished from other inflationary models of the same type by the next generation CMB satellites.
false
[ "≲ g ≲", "non-minimal coupling", "ϕ", "cosmic inflation", "≲", "particle decays", "the observed dark matter abundance", "a dark matter candidate", "other inflationary models", "the observed curvature perturbation amplitude", "equilibrium", "another scalar field", "the scalar field", "Universe", "the non-minimal coupling", "O}(1", "the next generation CMB satellites", "a Z <SUB>2</SUB>-symmetric scalar field", "ξ =", "CMB" ]
9.755808
-1.176382
77
2302031
[ "Biggs, A. D.", "Zwaan, M. A.", "Hatziminaoglou, E.", "Péroux, C.", "Liske, J." ]
2016MNRAS.462.2819B
[ "Parsec-scale H I absorption structure in a low-redshift galaxy seen against a compact symmetric object" ]
12
[ "European Southern Observatory, Karl Schwarzschild Straße 2, D-85748 Garching, Germany", "European Southern Observatory, Karl Schwarzschild Straße 2, D-85748 Garching, Germany", "European Southern Observatory, Karl Schwarzschild Straße 2, D-85748 Garching, Germany", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, F-13388 Marseille, France", "European Southern Observatory, Karl Schwarzschild Straße 2, D-85748 Garching, Germany; Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112, D-21029 Hamburg, Germany" ]
[ "2017MNRAS.465..588D", "2017MNRAS.465.4249D", "2017MNRAS.465.4450A", "2018MNRAS.476.2432G", "2018MNRAS.479L..50P", "2019JApA...40...41D", "2019MNRAS.482.2934A", "2022A&A...657A.113L", "2022JApA...43..103D", "2023MNRAS.519..931W", "2024ApJ...961..240K", "2024ApJ...961..242R" ]
[ "astronomy" ]
6
[ "galaxies: active", "galaxies: individual: J0855+5751", "galaxies: individual: SDSS J085519.05+575140.7", "galaxies: ISM", "radio lines: galaxies", "Astrophysics - Astrophysics of Galaxies" ]
[ "1974MNRAS.167P..31F", "1977ApJ...218..148M", "1986ApJ...303..617L", "1989PASP..101.1032M", "1994A&AS..106..275B", "1994ApJ...432L..87W", "1997ApJ...475..479W", "1997MNRAS.284..126B", "1998AJ....115.1693C", "1998PASP..110..493O", "1999ApJ...521..103P", "2001ApJ...548..749E", "2002ApJ...576..870P", "2002ApJS..141...13B", "2003A&A...404..871P", "2003ApJ...586.1067H", "2003ghr..conf..131B", "2005A&A...444L...9M", "2005ARA&A..43..861W", "2005ApJS..159...27T", "2005MNRAS.364.1467Z", "2006AJ....131.1163S", "2006MNRAS.371..431C", "2006MNRAS.372.1607T", "2006MNRAS.373..972G", "2006PASP..118.1711W", "2007ApJ...658..203H", "2007MNRAS.382.1639H", "2008ApJ...685..147N", "2008MNRAS.386.1252H", "2008MNRAS.387..639H", "2009MNRAS.399.1206H", "2010MNRAS.404L..45R", "2010evn..confE..86K", "2011MNRAS.413.1165C", "2011MNRAS.418.1787C", "2012A&A...547L...1N", "2012ApJ...754...29Z", "2012MNRAS.426..120F", "2013MNRAS.428.2198S", "2014ApJS..211...17A", "2014MNRAS.438.2131K", "2015MNRAS.453.1268Z" ]
[ "10.1093/mnras/stw1786", "10.48550/arXiv.1607.05995" ]
1607
1607.05995_arXiv.txt
\label{introduction} Much of what we know about the gas composition of galaxies has been made possible through observations of atomic hydrogen in absorption. In contrast to emission lines, the minimum optical depth that can be detected is set by the brightness of the background source and thus detections are almost independent of redshift. Detection of Ly-$\alpha$ absorption at optical and ultraviolet wavelengths has been very successful at providing information on the physical state of a wide variety of media, from the low-density intergalactic medium to high-redshift galaxies \citep*[e.g.][]{bechtold03,wolfe05,noterdaeme12}. A complementary probe is the \ion{H}{i} hyperfine 21-cm line, this having the additional advantage that it does not saturate and is unaffected by dust obscuration. The 21-cm line enables a determination of the kinematics and gas distribution in the intervening absorbers, which should then allow these systems to be linked to their $z=0$ galaxy counterparts seen in 21-cm emission \citep{zwaan05a}. One application of 21-cm absorption line studies that has been poorly exploited is the study of small-scale structure in the interstellar medium (ISM) of external galaxies. Observations of background structures on the order of tens of mas will probe the ISM of foreground galaxies at a distance of $\sim$100~Mpc on scales of tens of pc. This is interesting because it bridges the gap between the maximum of a few pc that can be studied in the ISM of the Milky Way in absorption against background radio sources \citep[e.g.][]{roy10}, and the more than 100-pc scales that are probed by high resolution 21-cm emission line maps of nearby galaxies \citep*[e.g.][]{elmegreen01,zhang12}. Neutral hydrogen opacity fluctuations on parsec scales provide us with information on the processes that regulate star formation. The structure and turbulence of the neutral medium eventually determine the size distribution of molecular clouds and affects the shape of the stellar initial mass function \citep{padoan02}. Therefore, measuring the small-scale structure of the ISM is essential for our understanding of how cold gas is converted into stars over cosmic time. \citet*{mckee77} modelled the ISM as a collection of small (0.4-10~pc) isotropic clouds, each containing a central cold core and a warmer envelope -- these two components constitute the cold and warm neutral media (CNM and WNM). More recently, an extensive analysis of Galactic 21-cm absorption lines led \citet{heiles03} to conclude that the ISM can be modelled by ``blobby sheets'' i.e.\ the CNM consists of sheet-like structures with blobs (cloudlets) embedded within. As an example of the study of an individual background source, \citet{srianand13} have presented a case of a quasar sight line piercing the gas disk of a galaxy at a redshift of $z=0.08$. Very Long Baseline Array (VLBA) observations resulted in 21-cm absorption spectra toward three components separated by $\sim$10 and 90~pc at the distance of the galaxy. The measured optical depths toward the two components differ by up to a factor of 10 for the narrowest components and much less for the broader component. This is interpreted as being due to small ($<10$~pc) dense clouds embedded in a diffuse neutral medium on scales of several tens of parsecs, in support of the \citeauthor{heiles03} model. \begin{figure*} \begin{center} \includegraphics[scale=0.15]{fig1a.png} \includegraphics[scale=0.15]{fig1b.png} \caption{Images of J0855+5751 at 2.3 (left) and 5~GHz (right) made from archival VLBA data. The 2.3-GHz data are taken from the VLBA Calibrator Survey \citep{beasley02} and at 5~GHz from VIPS \citep{taylor05,helmboldt07}. Our new analysis of these data phase-referenced both images to the position of the 8-arcmin distant source J0854+5757. The restoring beams are shown in the bottom-left corner and have dimensions of $8.1 \times 2.9$~mas$^2$ at 11\fdg0 (2.3~GHz) and $3.0 \times 1.4$~mas$^2$ at $-$1\fdg4 (5~GHz). Both maps are centred at 08$^{\rmn{h}}$~55$^{\rmn{m}}$~21\fs357, $+$57$\degr$~51\arcmin~44\farcs10 (J2000). Contours are plotted at multiples ($-1$, 1, 2, 4, 8, 16, etc.) of 3$\sigma$ where $\sigma$ is the off-source rms noise in the map ($\sigma_{2.3} = 2.2$~mJy\,beam$^{-1}$, $\sigma_{5} = 230 \, \mu$Jy\,beam$^{-1}$). The greyscale shows source brightness in mJy\,beam$^{-1}$. The source is very likely a CSO \citep{taylor05} with the faint central component (visible at 5~GHz only) corresponding to the location of the radio core.} \label{fig:vlba} \end{center} \end{figure*} J0855+5751 is a similar system that has been discovered in the course of a Green Bank Telescope (GBT) survey for small impact parameter pairs of optical galaxies and background radio-loud quasars \citep{zwaan15}. The dwarf galaxy SDSS~J085519.05+575140.7 at a redshift of $z=0.026$ displays strong 21-cm absorption with a peak optical depth of 24~per~cent. The background source, J0855+5751, has a 1.4-GHz flux density of 636~mJy \citep{white97} and a projected distance from the foreground galaxy of 6.8~kpc, a typical impact parameter for \ion{H}{i} absorption detected in the gas disk of a galaxy with a luminosity of 0.05~$L^*$ \citep{zwaan05a}. VLBA observations at 2.3 \citep{beasley02} and 5~GHz \citep{helmboldt07} have demonstrated that this source consists of two resolved components separated by about 60~mas (Fig.~\ref{fig:vlba}). On the basis of the observed VLBI structure, \citet{taylor05} have classified this source as a Compact Symmetric Object (CSO). Here we present global VLBI observations of J0855+5751 that have been used to map the \ion{H}{i} absorption in the foreground galaxy at high angular resolution, $\sim$4~mas, thus probing the spatial distribution of cold gas on scales ranging from 2 to 35~pc. In addition, we report on spectroscopy carried out with the William Herschel Telescope (WHT) with which we have measured the redshift of the background radio source's host galaxy. Throughout this paper we assume a flat $\Lambda$CDM cosmology and perform calculations using Ned Wright's Javascript Cosmology Calculator \citep{wright06}.
\label{conclusions} We have presented global VLBI observations of the bright radio source J0855+5751 which have allowed us to probe the properties of the \ion{H}{i} gas in the foreground dwarf galaxy SDSS~J085519.05+575140.7 at $z = 0.026$. We detect 21-cm absorption on all sightlines towards the background source where the SNR is high enough and, with a maximum velocity shift of $\le$2\kms, we seem to be probing a single coherent CNM structure with a minimum extent along one axis of 35~pc. The large size of this structure provides support for the applicability of the \citet{heiles03} ``blobby sheet'' model to the ISM of external galaxies. Our very high angular resolution ($3-4$~mas) has allowed us to probe and map the variation in the 21-cm optical depth against the two lobes of the background radio source on linear scales as small as $2$~pc. Whilst the absorption properties of the northern gas are rather uniform, those of the southern gas are much less so, with the total optical depth varying by a factor of five across a distance of $\sim$6~pc. The upper limit on the \ion{H}{i} column density lies significantly above the threshold that defines a DLA. Our observation of significant optical-depth variations on scales of order 2~pc suggests that $T_s$ measurements from low-redshift ($z < 2$) DLAs may be significantly more unreliable than reported by \citet{kanekar14}. The continuum map of the background radio source is by far the most sensitive yet made of J0855+5751 and detects much more of the extended emission in the radio lobes. The map includes a faint component previously seen at 5~GHz which, by virtue of its location and relatively flat spectrum, is probably the radio core. The new map thus strengthens the previous identification of this source as a Compact Symmetric Object. The southern lobe is distorted in comparison to its northern counterpart, most probably due to a collision with a dense cloud in the host galaxy. Given the CSO identification, future VLBI observations at 5~GHz might detect expansion of the source and lead to a measurement of a kinematic age. The redshift of the host galaxy has been measured using optical spectroscopy and is equal to $z = 0.54186 \pm 0.00009$. This offers the possibility of searching for absorption associated with the radio source. Associated \ion{H}{i} absorption is often seen in CSOs and other compact radio sources and, with the redshift of the host galaxy now known, a search for this could be made in J0855+5751. In addition, the host galaxy's very red colour ($V - K \gtrsim 5$) raises the possibility of detecting OH absorption in this system. Unfortunately, RFI and a lack of VLBI arrays with suitable receivers make such searches at redshifts greater than $\sim$0.1 difficult. This situation will change with the advent of the Square Kilometre Array (SKA) as this will be based in very benign RFI environments, be capable of tuning to effectively any plausible redshifted frequency of \ion{H}{i} or OH and also have baselines up to 3000~km \citep{morganti15}.
16
7
1607.05995
We present global VLBI observations of the 21-cm transition of atomic hydrogen seen in absorption against the radio source J0855+5751. The foreground absorber (SDSS J085519.05+575140.7) is a dwarf galaxy at z = 0.026. As the background source is heavily resolved by VLBI, the data allow us to map the properties of the foreground H I gas with a spatial resolution of 2 pc. The absorbing gas corresponds to a single coherent structure with an extent &gt;35 pc, but we also detect significant and coherent variations, including a change in the H I optical depth by a factor of 5 across a distance of ≲ 6 pc. The large size of the structure provides support for the Heiles &amp; Troland model of the interstellar medium, as well as its applicability to external galaxies. The large variations in H I optical depth also suggest that caution should be applied when interpreting T<SUB>S</SUB> measurements from radio-detected DLAs. In addition, the distorted appearance of the background radio source is indicative of a strong jet-cloud interaction in its host galaxy. We have measured its redshift (z = 0.541 86) using optical spectroscopy on the William Herschel Telescope and this confirms that J0855+5751 is an FR II radio source with a physical extent of &lt;1 kpc and supports the previous identification of this source as a compact symmetric object. These sources often show absorption associated with the host galaxy and we suggest that both H I and OH should be searched for in J0855+5751.
false
[ "H I optical depth", "J0855", "I optical depth", "external galaxies", "pc", "H", "optical spectroscopy", "William Herschel Telescope", "significant and coherent variations", "radio-detected DLAs", "OH", "z", "an FR II radio source", "Troland model", "the background radio source", "its host galaxy", "the host galaxy", "≲ 6 pc", "atomic hydrogen", "the radio source" ]
13.904736
8.249867
124
12470603
[ "Bret, Antoine" ]
2016JPlPh..82d9003B
[ "Particle trajectories in Weibel magnetic filaments with a flow-aligned magnetic field" ]
7
[ "ETSI Industriales, Universidad de Castilla-La Mancha, 13071 Ciudad Real, Spain" ]
[ "2017JPlPh..83b7101B", "2017LPB....35..513B", "2018JPlPh..84f9004B", "2019ApJ...876....2Y", "2019PhPl...26f2108B", "2020LPB....38..114B", "2021NJPh...23f3054J" ]
[ "astronomy", "physics" ]
2
[ "Physics - Plasma Physics", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1966RvPP....4...23S", "1972PhFl...15..317D", "1975PhFl...18..346G", "1987PhR...154....1B", "1998clel.book.....J", "1999ApJ...526..697M", "2002PhPl....9.2458S", "2003ApJ...592..378V", "2003ApJ...596L.121S", "2003ApJ...599L..57S", "2004RvMP...76.1143P", "2006ApJ...637..765M", "2006ApJ...647.1250L", "2006ApJ...651..584S", "2008PhRvL.100t5008B", "2009A&ARv..17..409T", "2011PhRvL.107b5003S", "2012ApJ...759...73N", "2012PhRvL.108w5004F", "2014NatSR...4E3934S", "2015NatPh..11..173H", "2015PhPl...22g2116B", "2016RPPh...79d6901M" ]
[ "10.1017/S0022377816000702", "10.48550/arXiv.1607.07442" ]
1607
1607.07442_arXiv.txt
Magnetic filaments are spontaneously generated by the growth of the Weibel, or filamentation, instability. This instability is triggered when two relativistic plasma shells cross each other \cite{Huntington2015}. It has been invoked for intergalactic magnetic fields generation \cite{Schlickeiser2003}, inertial confinement fusion \cite{Silva2002,DeutschPRE2005}, or collisionless shocks physics \cite{Medvedev1999,BretPoP2013,Fiuza2012}. Regarding collisionless shocks, it has been known for long that they are capable of accelerating particles \cite{Blandford1987} (for a particle to receive energy, and keep it, the medium has to be collisionless). As such, it is believed that they may play a key-role in the generation of high energy cosmic rays, or gamma ray bursts \cite{Vietri2003ApJ,Piran2005}. During the last decade or so, the physics of collisionless shocks has undergone a renewed interest, as it became possible to study them in detail through large-scale particle-in-cell simulations \cite{SilvaApJ,Spitkovsky2005,niemiec2012}. On the theory side, one of the problems currently being solved has to do with the very formation of such shocks. Concerning fluid shocks, it is known they can arise from the steepening of a large amplitude sound wave, or from the collision of two media, with a collision speed faster than the speed of sound in one of them \cite{Zeldovich}. The formation process in the collisionless case is more involved because such shocks occur in collionless plasmas. In these kind of settings, where the collision frequency is virtually zero, a shock has to be mediated by collective plasma interactions \cite{Sagdeev66}. As electromagnetic objects, collisionless shocks display far more variety than fluid shocks. They can be electrostatic or electromagnetic (Weibel) \cite{Sarri2011,Dieckmann2014NJP,Stockem2014}, form in pair or electron/ion plasmas \cite{BretPoP2013,BretPoP2014,Bret2015ApJL}, and on the top of these dichotomies, be influenced by the strength of an external magnetic field and its orientation \cite{Treumann2009,Marco2016}. Among the different kinds of collisionless shocks, an interesting sub-class is formed by the so-called ``Weibel shocks''. When two collisionless plasma shells run into each other, the overlapping region turns unstable. If the encounter occurs at relativistic velocity, the dominant instability is the Weibel one \cite{BretPoPHierarchie,BretPRL2008,BretPoPReview}. This instability grows magnetic filaments which can block the incoming flow, initiating the shock formation \cite{BretPoP2013,BretPoP2014,Bret2015ApJL}. Note that at the very beginning of the shock formation process, the counter-streaming plasmas cross each other, triggering the Weibel instability. At this stage, magnetic filaments are under formation, and particles are not stopped. Then, the Weibel instability reaches saturation, the filaments are fully formed, and may stop the particles arriving at later times. The flow stopped by the filaments is not the flow that formed them. It is clear that in shock setting, once the filaments have stopped a number of particles, these may perturb the filaments and break them. Yet, it has been found that the resulting field configuration is at least as efficient as were the filaments to stop the flow \cite{BretPoP2014,Bret2015ApJL}. In order to know whether a shock will start forming after the saturation of the Weibel instability, or not, the key question is therefore: at saturation time, are the filaments able to block the incoming particles? For the case where there is no external magnetic field, the conditions upon which the flow is blocked can been derived from the analysis of the motion of charged particles in the Weibel magnetic filaments \cite{BretPoP2015}. Yet, many astrophysical settings are magnetized. If two plasma shells interact over an external, flow-aligned, magnetic field $\mathbf{B}_0$, the Weibel instability still grows while the field is not too strong \cite{Godfrey1975,BretPoPMagne}. Here, the resulting magnetic filaments supposed to block the incoming flow for the shock to form, will be superimposed over the external $\mathbf{B}_0$. The goal of this article is therefore to study this problem: how about the trajectory of charged particles within magnetic filaments and an external $\mathbf{B}_0$?
The formation of a Weibel shock involves to collision of two plasma shells. As they pass through each other, the overlapping region becomes Weibel unstable. At saturation, it forms magnetic filaments of transverse size $1/k$ and peak field $B_f$. We here determined under which conditions these filaments block the incoming flow, thus initiating the shock formation. It has already been found than in the absence of an external magnetic field, most of the flow (if cold) keeps streaming through the filaments, or not, whether $1/k$ is smaller or larger than the Larmor radius of the particles in the field $B_f$. In the presence of a parallel magnetic field $\mathbf{B}_0$, everything relies on the parameter $\alpha=B_0/B_f$. While $\alpha \lesssim 0.4$, the overall picture is similar to that with $\alpha =0$. But for $\alpha>1/2$ a transition occurs. The guiding field $\mathbf{B}_0$ becomes dominant and particles keep streaming through the filaments regardless of their initial velocities. In realistic scenarios however, the parameter $\sigma$ which governs the transition, forces $\alpha \ll 1$. Further works could contemplate the case of a perpendicular field. Such a scenario involves more free parameters than the current one since the field will have to be specified by its two normal components, instead of one single parallel component in the present case.
16
7
1607.07442
For a Weibel shock to form, two plasma shells have to collide and trigger the Weibel instability. At saturation, this instability generates magnetic filaments in the overlapping region with peak field f]]&gt; . In the absence of an external guiding magnetic field, these filaments can block the incoming flow, initiating the shock formation, if their size is larger than the Larmor radius of the incoming particles in the peak field. Here we show that this result still holds in the presence of an external magnetic field, provided it is not too high. Yet, for 0\gtrsim B<SUB>f</SUB>/2, the filaments become unable to stop any particle, regardless of its initial velocity.
false
[ "peak field", "magnetic filaments", "an external guiding magnetic field", "f]]&gt", "an external magnetic field", "the peak field", "Weibel", "Larmor", "the incoming particles", "the incoming flow", "its initial velocity", "a Weibel shock", "the Weibel instability", "the shock formation", "the Larmor radius", "two plasma shells", "the filaments", "these filaments", "any particle", "the overlapping region" ]
7.278462
4.339656
31
12473525
[ "Drozdovskaya, Maria N.", "Walsh, Catherine", "van Dishoeck, Ewine F.", "Furuya, Kenji", "Marboeuf, Ulysse", "Thiabaud, Amaury", "Harsono, Daniel", "Visser, Ruud" ]
2016MNRAS.462..977D
[ "Cometary ices in forming protoplanetary disc midplanes" ]
82
[ "Leiden Observatory, Leiden University, PO Box 9513, NL-2300 RA Leiden, the Netherlands", "Leiden Observatory, Leiden University, PO Box 9513, NL-2300 RA Leiden, the Netherlands", "Leiden Observatory, Leiden University, PO Box 9513, NL-2300 RA Leiden, the Netherlands; Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany", "Leiden Observatory, Leiden University, PO Box 9513, NL-2300 RA Leiden, the Netherlands", "Center for Space and Habitability, Universität Bern, CH-3012 Bern, Switzerland; NCCR PlanetS-Universität Bern, Physikalisches Institut, Universität Bern, CH-3012 Bern, Switzerland", "Center for Space and Habitability, Universität Bern, CH-3012 Bern, Switzerland; NCCR PlanetS-Universität Bern, Physikalisches Institut, Universität Bern, CH-3012 Bern, Switzerland", "Center for Astronomy, Institute for Theoretical Astrophysics, Heidelberg University, Albert-Ueberle-Strasse 2, D-69120 Heidelberg, Germany", "European Southern Observatory, Karl-Schwarzschild-Strasse 2, D-85748 Garching, Germany" ]
[ "2016A&A...596L...3I", "2016ApJ...833..105Y", "2016PASA...33...53H", "2017A&A...599A..40F", "2017A&A...599A.132L", "2017A&A...601A..36B", "2017AJ....153..168R", "2017AJ....153..241B", "2017ApJ...837...78H", "2017ApJ...837..177C", "2017ApJ...841...39Y", "2017MNRAS.467.2552C", "2017MNRAS.467.4753N", "2017MNRAS.469S.818L", "2017RSPTA.37560252B", "2018A&A...611A..80B", "2018A&A...612A..88L", "2018A&A...613A..14E", "2018A&A...615A..83V", "2018A&A...618A.182B", "2018ARep...62..455S", "2018ASPC..517....3M", "2018ASPC..517...73C", "2018ApJ...866...46M", "2018ApJ...868....9Q", "2018ApJ...868L..37D", "2018IAUS..332....3V", "2018IAUS..332..395W", "2018MNRAS.475.2355M", "2018arXiv181007867C", "2019A&A...627A.127C", "2019A&A...631A..25J", "2019AJ....158..128M", "2019ApJ...885..146R", "2019ApJ...886....6T", "2019ESC.....3.1792R", "2019ESC.....3.2158E", "2019MNRAS.484..345C", "2019MNRAS.484.1563W", "2019MNRAS.485.1843W", "2019MNRAS.490...50D", "2019arXiv190711081B", "2020A&A...635A..48M", "2020A&A...639A..87V", "2020A&A...643A.108C", "2020AJ....159...42R", "2020ApJ...905..162T", "2020IAUS..345..115K", "2020IAUS..350..207K", "2020MNRAS.498..276K", "2020NatAs...4..861C", "2020Sci...367.3705G", "2021A&A...648A..24V", "2021A&A...650A.180N", "2021A&A...656A.146M", "2021ApJ...908..108F", "2021ApJ...919...45B", "2021ApJS..257....9I", "2021LPICo2609.6173B", "2021MNRAS.500.4658M", "2021PSJ.....2...45B", "2021PhR...893....1O", "2022A&A...666A..35T", "2022ApJ...931..164W", "2022ApJ...936..188N", "2022ApJ...938...29F", "2022ESC.....6.1171K", "2022ExA....54.1077B", "2022Life...12..404P", "2022MNRAS.510.1148S", "2022PASA...39....9A", "2022arXiv221016354F", "2022arXiv221214529A", "2023A&A...670A.141T", "2023A&A...678A..33J", "2023A&A...679A.138S", "2023ARA&A..61..287O", "2023ApJ...942...50P", "2023ApJ...954..167P", "2023ApJ...955....5S", "2024A&A...685A.112N", "2024arXiv240602387A" ]
[ "astronomy" ]
11
[ "astrochemistry", "comets: general", "protoplanetary discs", "stars: protostars", "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1093/mnras/stw1632", "10.48550/arXiv.1607.07861" ]
1607
1607.07861.txt
\label{intro} Protoplanetary discs encircling young protostars are the birth locations of future mature planetary systems (see \citealt{Johansen2014} for a review). Past observations have suggested that dust growth from micron sizes may begin as early as the prestellar core stage and dust grains may reach millimetre dimensions during infall in the envelopes of protostars (\citealt{Pagani2010, Miotello2014, Jones2016, Ysard2016}, see \citealt{Testi2014} for a review). Recent ALMA data have shown strongly defined ring structures in the disc around the Class I--II protostar HL Tau, which is thought to be younger than $\leq 1 - 2$~Myr and still embedded in a large envelope \citep{ALMA2015}. Models have suggested that these rings may be caused by several planets of at least $0.2$~M$_{\text{J}}$ in mass clearing gaps in the dust distribution \citep{Dipierro2015, Pinte2016, Dong2015}. An alternative hypothesis is that the contrast reflects enhanced dust growth to centimeter sizes behind various snowlines, leading to a change in opacity in the emitting dust \citep{KeZhang2015}. This is potentially aided by sintering close to the snowlines of the volatiles \citep{Okuzumi2016}. These findings are pushing the onset of planetesimal formation much earlier along the star-disc evolutionary sequence than previously thought, perhaps even as early as the embedded phase. At the same time, the exoplanet community has demonstrated the diverse outcomes of planet formation (see, e.g., www.exoplanet.eu, \citealt{Schneider2011}). Planet population synthesis modellers have carried out pioneering work in linking protoplanetary disc theory with the final architecture of planetary systems. Such models take the key physical processes across all these evolutionary stages into account, and with large sets of initial conditions, to make statistical predictions on the exoplanet population (see \citealt{Benz2014} for a review). It has been postulated that planetary atmospheres form initially via pebble accretion and heating from this assembly prior to runaway gas accretion \citep{InabaIkoma2003, OrmelKlahr2010, Bitsch2015}. Based on the newest results from the star formation community, it may be necessary to use the initial conditions from the earlier embedded phase rather than the classical T Tauri (Class II) stage. Typically, the constituents of planetesimals that feed protoplanets are grouped into volatiles (H, O, C, N and S containing molecules) and refractories/rocks (minerals/inorganic refractories and complex organic refractory components such as PAHs and macromolecular complex organic matter). For regions of a protoplanetary disc where most volatiles are frozen out as solids, an ice/rock mass ratio of $\sim 2 - 4$ is suggested \citep{Pontoppidan2014}. Various instruments aboard the \textit{Rosetta} mission are attempting to constrain the the icy and dusty contents of comet 67P/Churyumov-–Gerasimenko to an unprecedented precision. The RSI experiment indicates a ice/dust mass of ratio of $\sim 4$ based on the gravity field of the comet \citep{Patzold2016}. CONSERT measurements of an ice/dust volumetric ratio of $\sim 0.4-2.5$ \citep{Kofman2015} imply a ice/dust mass ratio of $\sim 0.1-0.9$, assuming an average dust density of $2~600$~kg~m$^{-3}$ and an ice density of $940$~kg~m$^{-3}$ as in \citet{Patzold2016}. These ratios mean that the solid-phase chemical composition is a pivotal parameter for the subsequent protoplanet makeup. The density is highest in the midplane of a protoplanetary disc, thus harbouring the bulk of the mass and implying that predominantly the ices in that region will likely shape the chemical composition of the atmospheres of the exoplanets that are observed today. The chemical composition of discs in the embedded phase has been probed by models of \citet{Visser2009}, \citet{Visser2011}, \citet{Drozdovskaya2014}, \citet{Harsono2015}. Observational evidence for them is scarce, because high spatial resolution ($\lesssim 30$~AU) is necessary to get spatial information for a small ($\sim 100$~AU) target,whose emission is easily overwhelmed by that of the more massive envelope, making it hard to constrain their physical parameters. Only recently, observations have shown embedded Keplerian discs of $\sim 50 - 300$~AU in radius in embedded protostars \citep{Tobin2012, Tobin2013, Tobin2015, Brinch2013, Murillo2013, Sakai2014, Harsono2014, Chou2014}. Instead, significant modelling effort has focused on discs in the Class II/III stages of star formation (see, e.g., \citet{HenningSemenov2013} and table~$1$ in \citealt{Woitke2016} and the references therein). The chemistry in such Class II discs has been modelled by means of gas-grain models with grain-surface reactions, e.g., \citet{Willacy2006, Semenov2006, Hersant2009, Semenov2010, Henning2010, Walsh2010, Walsh2012, Vasyunin2011, Akimkin2013}. Only recently complex organics (loosely defined in both chemistry and astronomy as large, $\geq 6$ atoms, carbon-containing species, \citealt{HerbstvanDishoeck2009}) have also been considered by \citet{SemenovWiebe2011}, \citet{Furuya2014}, \citet{Walsh2014a}, \citet{Walsh2014b}, see also \citet{HenningSemenov2013} for a review. Meanwhile, such molecules have been frequently observed towards protostars and in Solar System bodies \citep{HerbstvanDishoeck2009}. The aim of this work is to compute the chemical composition of planetesimals and cometary materials in the midplanes of protoplanetary discs in the embedded phase of star formation. This is done using state-of-the-art physicochemical models that are dynamic in nature and include chemical kinetics, building upon the work in \citet{Drozdovskaya2014}, which studied the discs as a whole. Main volatiles that are expected to dominate the icy component of plantesimals in the midplanes of protoplanetary discs are considered first. Trace complex organic species are also analysed in this work, since they are potentially the earliest precursors to prebiotic molecules. The dynamic nature of the model allows the study of the chemical effects stemming from the transport from the prestellar cloud through the envelope and into the disc. The key model features are highlighted in Section~\ref{model}. The results are shown in Section~\ref{results} and are discussed in Section~\ref{discussion}, including all the derived implications for the population synthesis models and the Solar System community. The conclusions are presented in Section~\ref{conclusions}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\label{conclusions} This paper builds on previous work by \citet{Visser2009, Visser2011} and \citet{Drozdovskaya2014} and is focused on the midplane composition of protoplanetary discs in the embedded phase. State-of-the-art physicochemical models are employed to simulate the formation of discs from the initial prestellar phase for $2.46 \times 10^{5}$~yr. Two discs are studied that vary in their respective dominant disc-growth mechanism, either viscous spreading or pure infall. Subsequently, the path of material towards the midplanes of these discs differs, predominantly in-out or simply inwards, respectively. These routes in turn set the temperatures and UV fluxes during the transport of parcels from cloud to disc. More than $100$ trajectories into each disc are computed to sample the disc midplane, and the icy content in the framework of such a dynamical model is analysed thereafter. The main conclusions are as follows: \begin{itemize} \item The typical main ice constituents in the midplanes are $\sim 70-80$\% H$_{2}$O, $\sim 10-20$\% CO$_{2}$, $1-10$\% NH$_{3}$ and $1-4$\% CH$_{3}$OH; however, CO$_{2}$ may dominate ($\sim 40$\%) in the outer disc when grown via pure infall. \item Trace complex organic ices are most abundant in the outer disc ($R \gtrsim 30$~AU). Some may be as plentiful as methanol ice in the midplane, at a level of $\sim 1$\% of the total ice content, similar to that seen in comet observations. The inner disc is rich in complex organic gases. \item The positions of snowlines of volatiles and complex organics in the midplanes of protoplanetary discs, according to relative volatilities, are retained even when dynamics and chemistry during disc formation are taken into account. \item Dynamic infall and the chemistry en route to the midplane may enhance the amount of CO$_{2}$ and diminish CH$_{3}$OH ices in comparison to the prestellar phase. Not all volatiles are simply inherited by the midplane from the cloud. Icy planetesimals and cometary bodies reflect the provenances of the midplane ices. \item The elevated temperatures and additional FUV photons in the envelope facilitate the formation of prebiotically-significant molecules, which may consitute as much as $\sim 10$\% of the icy mantles. Current Class II disc models may be underestimating the complex organic content, since the initial abundances set in the embedded phase could be much higher than currently assumed. \item The C/O and C/N ratios differ between the gas and solid phases. The two ratios in the ice show little variation beyond the inner $10$~AU and both are nearly solar in the case of pure infall. %\item The typical main ice constituents in the midplanes are $\sim 70-80$\% H$_{2}$O, $\sim 10-20$\% CO$_{2}$, $1-10$\% NH$_{3}$ and $1-4$\% CH$_{3}$OH, however CO$_{2}$ may dominate in the outer disc when grown via pure infall. %\item Trace complex organic ices are most abundant in the outer disc ($R \gtrsim 30$~AU). Some may be as plentiful as methanol ice in the midplane, at a level of $\sim 1$\% of the total ice content, similar to that seen in comet observations. %\item The positions of snowlines of volatiles and complex organics in the midplanes of protoplanetary discs, according to relative volatilities, are retained even when dynamics and chemistry during disc formation are taken into account. %\item An inner reservoir of gaseous complex organics appears to originate from the inner few AU of the midplane that is sufficiently hot to thermal desorb such tightly bound molecules and shielded enough to prevent their photodissociation by intense protostellar FUV photons in such proximity to the protostar. %\item Dynamic infall is crucial in facilitating complex organic molecule formation thanks to elevated temperatures and additional FUV photons en route to the disc midplane. %\item Planet population synthesis modellers may be underestimating the amount of CO$_{2}$ and overestimating CH$_{3}$OH ices in planetesimals by neglecting chemical processing, which can affect the C/O ratio in planetesimals. By excluding complex organics from their models, they may be missing as much as $\sim 10$\% of the ice available in midplanes for planetesimals, which are additionally prebiotically important. %\item Volatiles are not simply inherited from the prestellar phase by the disc midplanes and chemistry en route does modify the composition of the icy mantle. %\item The prebiotically-significant content of icy planetesimals and cometary bodies is determined by the chemistry in the envelope en route, before the material even enters the protoplanetary disc. %\item The chemical complexity of comets can be explained by the complex organic ice budget in the outer regions of protosolar discs. \end{itemize} The latest theories and observations are suggesting much earlier planetesimal formation than previously thought. The results presented in this paper for the midplanes of protoplanetary discs around low-mass protostars in the embedded phase may probe the volatile and prebiotically-significant content of the pebbles that go on to feed exoplanet atmospheres and form comets. Future work will focus on deeper understanding of the consequences of these model results on cometary composition and planetary populations. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
16
7
1607.07861
Low-mass protostars are the extrasolar analogues of the natal Solar system. Sophisticated physicochemical models are used to simulate the formation of two protoplanetary discs from the initial prestellar phase, one dominated by viscous spreading and the other by pure infall. The results show that the volatile prestellar fingerprint is modified by the chemistry en route into the disc. This holds relatively independent of initial abundances and chemical parameters: physical conditions are more important. The amount of CO<SUB>2</SUB> increases via the grain-surface reaction of OH with CO, which is enhanced by photodissociation of H<SUB>2</SUB>O ice. Complex organic molecules are produced during transport through the envelope at the expense of CH<SUB>3</SUB>OH ice. Their abundances can be comparable to that of methanol ice (few per cent of water ice) at large disc radii (R &gt; 30 au). Current Class II disc models may be underestimating the complex organic content. Planet population synthesis models may underestimate the amount of CO<SUB>2</SUB> and overestimate CH<SUB>3</SUB>OH ices in planetesimals by disregarding chemical processing between the cloud and disc phases. The overall C/O and C/N ratios differ between the gas and solid phases. The two ice ratios show little variation beyond the inner 10 au and both are nearly solar in the case of pure infall, but both are subsolar when viscous spreading dominates. Chemistry in the protostellar envelope en route to the protoplanetary disc sets the initial volatile and prebiotically significant content of icy planetesimals and cometary bodies. Comets are thus potentially reflecting the provenances of the midplane ices in the solar nebula.
false
[ "pure infall", "viscous spreading", "large disc radii", "water ice", "dominates", "Current Class II disc models", "cometary bodies", "icy planetesimals", "overestimate CH<SUB>3</SUB>OH ices", "chemical processing", "CH<SUB>3</SUB>OH ice", "initial abundances", "chemical parameters", "planetesimals", "the initial prestellar phase", "Complex organic molecules", "physical conditions", "OH", "little variation", "Planet population synthesis models" ]
9.414617
14.187407
-1
12408897
[ "Wood, Brian E.", "Müller, Hans-Reinhard", "Harper, Graham M." ]
2016ApJ...829...74W
[ "Hubble Space Telescope Constraints on the Winds and Astrospheres of Red Giant Stars" ]
15
[ "Naval Research Laboratory, Space Science Division, Washington, DC 20375, USA", "Department of Physics and Astronomy, Dartmouth College, Hanover, NH 03755, USA", "CASA, University of Colorado, Boulder, CO 80309-0389, USA" ]
[ "2018ApJ...869....1R", "2018ApJ...869..157C", "2018JPhCS1100a2028W", "2019ApJ...879...77Y", "2020A&A...635A..52M", "2021ApJ...915...37W", "2021ApJ...923...43V", "2021LRSP...18....3V", "2021MNRAS.500.3438O", "2021MNRAS.506.1697V", "2022A&A...666A.197C", "2022ApJ...932...57H", "2022MNRAS.517.6077E", "2023A&A...676A..98S", "2024ApJ...967..120W" ]
[ "astronomy" ]
8
[ "stars: chromospheres", "stars: late-type", "stars: winds", "outflows", "ultraviolet: stars", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1970PASP...82..169L", "1977A&A....57..395R", "1978ApJ...220..962K", "1979ApJ...229L..27L", "1979ApJ...231..128B", "1979ApJ...234.1023B", "1982A&A...107..292R", "1982A&A...115..280B", "1986AJ.....91..602D", "1986ESASP.263..193D", "1987A&A...172..111S", "1989A&A...216..139S", "1991ApJ...371..357J", "1991ApJ...383L..15H", "1992A&A...256..185C", "1992pavi.book.....S", "1994SoPh..152...69D", "1995A&A...300..775M", "1995ApJ...452..407H", "1995JGR...10021595P", "1996ApJ...463..254L", "1997A&A...318..215S", "1997ApJ...488..760D", "1997ApJ...491..876A", "1998ApJ...492L..83K", "1998ApJ...493..715S", "1998ApJ...494..700B", "1998ApJ...494..828J", "1998ApJ...503..396R", "1998MNRAS.298..387D", "1998PASP..110.1183W", "1998SSRv...85..161G", "1999A&A...346..785S", "1999A&AS..139..433D", "1999PASP..111.1515V", "2002ApJ...580..938J", "2002ApJS..139..439R", "2003A&A...411..447L", "2003AJ....126.2502M", "2003ApJ...594L..51R", "2003ApJ...598..610A", "2003ApJS..149..205M", "2003yCat.2246....0C", "2004A&A...424..727P", "2004ApJ...602..776R", "2005ApJ...622..680R", "2005ApJ...623L.137K", "2005ApJ...628L.143W", "2005ApJ...631L.155U", "2005ApJS..159..118W", "2006ApJ...651.1126A", "2007A&A...474..653V", "2007ApJ...655..946W", "2007ApJ...657L..57C", "2007ApJ...659.1592S", "2008ApJ...673..283R", "2010A&A...516L...2A", "2010ApJ...723.1210A", "2010MNRAS.403.1592B", "2011ASPC..448.1145H", "2011MNRAS.414..418P", "2012A&A...540A.130S", "2013AJ....146...98O", "2013ApJ...772...89L", "2014A&A...561A..85L", "2014A&A...561A.126D", "2014ApJ...780..108W", "2014ApJ...781L..33W", "2014ApJ...791L..14S", "2015ApJ...801...62W" ]
[ "10.3847/0004-637X/829/2/74", "10.48550/arXiv.1607.07732" ]
1607
1607.07732_arXiv.txt
Ultraviolet and X-ray observations have demonstrated that among main sequence stars stellar coronae are a ubiquitous phenomenon, with coronal emissions being detectable for all stars with spectral types later than A5~V \citep[e.g.,][]{js97}. This is not the case, however, above the main sequence. For giant stars, F and G type giants are coronal, but this coronal emission either disappears entirely or at least dramatically weakens for K and M type giants. For the K and M giants, the coronae seem to be largely replaced by strong winds ($\dot{M}\sim 10^{-11}$ M$_{\odot}$~yr$^{-1}$) at chromospheric temperatures ($T\sim 10^4$~K). The first recognition of the existence of this coronal dividing line among evolved stars between coronal and ``windy'' stars was by \citet{jll79}. If attention is limited to giant stars, i.e.\ luminosity class III stars, the dividing line lies at a spectral type of about K2~III \citep{bh91}. The nature of the transition across this dividing line is a topic of particular interest. Some stars to the right of the dividing line, referred to as ``hybrid chromosphere'' stars, seem to maintain some level of coronal emission, with $\gamma$~Dra (K5~III) being perhaps the best-studied luminosity class III example \citep{tra97}. This has led to the proposition that magnetic loops with coronal temperatures still persist beyond the dividing line, but the coronal emissions from these loops are completely or mostly hidden by the dense, thick chromospheres that have developed for such stars, from which emanate the strong chromospheric winds \citep{tra03}. In order to further investigate the transition from coronae to chromospheric winds, we here present a new survey of red giant stars with spectral types later than K2~III, using UV spectra taken by the {\em Hubble Space Telescope} (HST). We are not only interested in the winds of these stars, but also in the interactions of these winds with the interstellar medium (ISM). Ultraviolet spectroscopy from HST has provided the first spectroscopic detections of the wind-ISM interaction regions, or ``astrospheres,'' of many cool stars, including the heliospheric structure surrounding our own Sun. Hydrogen Lyman-$\alpha$ absorption has been observed from the ``hydrogen wall'' region outside the heliopause \citep{jll96}, and from the heliotail \citep{bew14a}. Hydrogen wall absorption has also been observed around other Sun-like stars, representing the first method by which solar-like coronal winds can be detected around other stars \citep{bew05a,bew14b}. However, we are here interested in the astrospheres of red giant stars. Observations of the Mg~II h \& k lines near 2800~\AA\ have been used to detect absorption from the astrosphere of $\alpha$~Tau (K5~III) \citep{rdr98,bew07}. Models of the $\alpha$~Tau astrosphere were computed using codes designed to study the heliosphere. Analogous to the solar wind, the red giant wind expands radially from the star until it reaches the termination shock (TS), where it is shocked to subsonic speeds. The post-TS wind is slower, hotter, and more dense. It is this post-TS region that provides Mg~II column densities sufficiently high to yield the observed Mg~II astrospheric absorption. Due to its high positive stellar radial velocity ($V_{rad}=+54.3$ km~s$^{-1}$), our line of sight towards $\alpha$~Tau is estimated to be at an angle of $\theta\sim 170^{\circ}$ from the upwind direction of the ISM flow vector in the rest frame of the star. This line of sight is therefore very much through the ``astrotail,'' and this may be essential for detecting the astrospheric absorption. A downwind line of sight provides a much longer path length through the post-TS material than an upwind or sidewind line of sight, leading to much higher column densities through the post-TS material \citep{bew07}. Nevertheless, the hydrodynamic models of the $\alpha$~Tau astrosphere were unable to reproduce the observed astrospheric absorption very well, predicting far too little absorption, and placing the absorption farther from the stellar rest frame than observed. The proposed explanation was that the $\alpha$~Tau TS is a radiative shock, with radiative cooling from H Lyman lines yielding further compression and deceleration behind the TS, resulting in a stronger Mg~II absorption feature closer to the rest frame of the star, in better agreement with the observed absorption \citep{bew07}. The $\alpha$~Tau observation has long been the only example of a detected red giant astrosphere. Our new HST survey of red giant stars was not only designed to study the red giant winds, but also designed to try to find new detections of astrospheric absorption around such stars.
We have analyzed the chromospheric Mg~II h \& k lines of K2-M5~III stars observed by HST, consisting of 9 observations obtained as part of a new HST red giant survey, and 4 archival targets. Our findings are summarized as follows: \begin{enumerate} \item The Mg~II line profiles of all 13 stars in our sample show evidence for stellar wind absorption, but for three of the K2-K4~III stars the effect is only an induced asymmetry in the line profile rather than a deep wind absorption feature. \item Measured Mg~II surface fluxes are very tightly correlated with spectral type and photospheric temperature, consistent with the idea that K2-M5~III stars redward of the coronal dividing line are all emitting at a basal flux level. The Mg~II k/h flux ratio increases towards later spectral types. \item Wind speeds estimated empirically from the Mg~II spectra correlate with spectral type and photospheric temperature, with $V_w$ decreasing from $V_w\approx 40$ km~s$^{-1}$ at K2~III to $V_w\approx 20$ km~s$^{-1}$ at M5~III. \item There are 2 new detections of astrospheric absorption among the recently observed stars ($\gamma$~Eri and $\sigma$~Pup), for a total of 3 in our sample, including the previous detection towards $\alpha$~Tau. However, the limited number of new detections indicates that detectable astrospheric absorption in UV lines is not a common phenomenon. For both $\gamma$~Eri and $\sigma$~Pup we detect astrospheric Fe~II $\lambda$2600.2 absorption in addition to the Mg~II signature, and for $\sigma$~Pup, astrospheric absorption is also observed in C~II $\lambda$1334.5. These are the first detections of red giant astrospheres in lines other than Mg~II. \item Analysis of the astrospheric absorption leads to measurements of TS compression ratio and post-TS temperature for the three detected astrospheres. The temperatures are correlated with $V_w$. However, the $T=(0.9-2.2)\times 10^4$~K post-TS temperatures are too low and the $\eta=6-18$ compression ratios too high according to the Rankine-Hugoniot shock jump conditions, providing further evidence that red giant termination shocks are radiative shocks rather than simple hydrodynamic shocks. \item We compute a hydrodynamic model of the $\gamma$~Eri astrosphere, which is the one with the TS compression ratio ($\eta=5.9\pm 1.7$) closest to the strong hydrodynamic shock value of $\eta=4$. Not surprisingly, this model has less difficulty reproducing the observed absorption than was the case for a past study of the $\alpha$~Tau astrosphere \citep{bew07}, for which $\eta=17.5\pm 3.1$. The $\gamma$~Eri model places the absorption at about the right velocity and with about the correct width, but it underpredicts the Mg~II opacity by a factor of 8. This might be due to either an underestimate of the Mg abundance in the stellar wind, or an underestimate of the ISM pressure surrounding the star. \end{enumerate}
16
7
1607.07732
We report on an ultraviolet spectroscopic survey of red giants observed by the Hubble Space Telescope, focusing on spectra of the Mg II h and k lines near 2800 Å in order to study stellar chromospheric emission, winds, and astrospheric absorption. We focus on spectral types between K2 III and M5 III, a spectral type range with stars that are noncoronal, but possessing strong, chromospheric winds. We find a very tight relation between Mg II surface flux and photospheric temperature, supporting the notion that all K2-M5 III stars are emitting at a basal flux level. Wind velocities (V <SUB> w </SUB>) are generally found to decrease with spectral type, with V <SUB> w </SUB> decreasing from ∼40 km s<SUP>-1</SUP> at K2 III to ∼20 km s<SUP>-1</SUP> at M5 III. We find two new detections of astrospheric absorption, for σ Pup (K5 III) and γ Eri (M1 III). This absorption signature had previously only been detected for α Tau (K5 III). For the three astrospheric detections, the temperature of the wind after the termination shock (TS) correlates with V <SUB> w </SUB>, but is lower than predicted by the Rankine-Hugoniot shock jump conditions, consistent with the idea that red giant TSs are radiative shocks rather than simple hydrodynamic shocks. A full hydrodynamic simulation of the γ Eri astrosphere is provided to explore this further. <P />Based on observations made with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with program GO-13462. This paper also presents observations obtained with the Harlan J. Smith Telescope at McDonald Observatory of the University of Texas at Austin.
false
[ "M5 III", "K2 III", "K5 III", "simple hydrodynamic shocks", "ESA Hubble Space Telescope", "shocks", "∼20 km s", "spectral type", "spectral types", "∼40 km s", "stellar chromospheric emission", "Wind velocities", "NASA contract NAS", "winds", "red giant TSs", "Mg II surface flux", "astrospheric absorption", "K2-M5 III", "SUB", "(K5 III" ]
8.876681
10.866913
153
12472962
[ "Rozehnal, J.", "Brož, M.", "Nesvorný, D.", "Durda, D. D.", "Walsh, K.", "Richardson, D. C.", "Asphaug, E." ]
2016MNRAS.462.2319R
[ "Hektor - an exceptional D-type family among Jovian Trojans" ]
18
[ "Institute of Astronomy, Charles University, Prague, V Holešovičkách 2, CZ-18000 Prague 8, Czech Republic; Štefánik Observatory, Petřín 205, CZ-11800 Prague, Czech Republic", "Institute of Astronomy, Charles University, Prague, V Holešovičkách 2, CZ-18000 Prague 8, Czech Republic", "Southwest Research Institute, 1050 Walnut St, Boulder, CO 80302, USA", "Southwest Research Institute, 1050 Walnut St, Boulder, CO 80302, USA", "Southwest Research Institute, 1050 Walnut St, Boulder, CO 80302, USA", "Department of Astronomy, University of Maryland, College Park, MD 20742-2421, USA", "School of Earth and Space Exploration, Arizona State University, Tempe, AZ 85287, USA" ]
[ "2017Icar..288..240M", "2018A&A...620A.167S", "2018ARA&A..56..137N", "2018AdSpR..61.1371J", "2018MNRAS.475..974P", "2019AJ....157..181V", "2020MNRAS.495.4085H", "2020MNRAS.499.3630H", "2021A&A...654A..75V", "2021MNRAS.504.1571H", "2021PSJ.....2..170B", "2021RNAAS...5...42B", "2022AJ....164..167M", "2022PSJ.....3..190H", "2023AJ....165...15W", "2023PSJ.....4..168B", "2023SSRv..219...83B", "2024AJ....167..138V" ]
[ "astronomy" ]
5
[ "celestial mechanics", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1962Natur.195..763T", "1969JGR....74.2531D", "1994AJ....107..772Z", "1994IAUS..160..205F", "1994Icar..108...18L", "1998A&A...339..272D", "1999Icar..142....5B", "2000Icar..143...45R", "2001CeMDA..80...39L", "2001Icar..153..348C", "2002SPIE.4836...98I", "2002aste.book...27B", "2004ApJ...614..497G", "2005Natur.435..462M", "2006Icar..182..496E", "2007Icar..186..498D", "2009A&A...493.1125L", "2009Icar..204..558M", "2009Natur.460..364L", "2010AJ....140.1391M", "2011AJ....141...25E", "2011ApJ...741...90M", "2011ApJ...742...40G", "2011MNRAS.414..565B", "2011PASJ...63.1117U", "2012ApJ...759...49G", "2013A&A...551A.117B", "2013ApJ...768...45N", "2013Icar..223..844B", "2014ApJ...783L..37M", "2015MNRAS.451.2399C", "2015MNRAS.454.2436V", "2015aste.book..297N", "2016MNRAS.457.1332C" ]
[ "10.1093/mnras/stw1719", "10.48550/arXiv.1607.04677" ]
1607
1607.04677_arXiv.txt
Jovian Trojans are actually large populations of minor bodies in the 1:1 mean motion resonance (MMR) with Jupiter, librating around $L_4$ and $L_5$ Lagrangian points. In general, there are two classes of theories explaining their origin: i) a theory in the framework of accretion model (e.g. Goldreich 2004, Lyra et al. 2009) and ii) a capture of bodies located in libration zones during a migration of giant planets (Morbidelli et al. 2005, Morbidelli et al. 2010, Nesvorn\'y et al. 2013), which is preferred in our solar system. Since the librating regions are very stable in the current configuration of planets and they are surrounded by strongly chaotic separatrices, bodies from other source regions (e.g. Main belt, Centaurs, Jupiter family comets) cannot otherwise enter the libration zones and Jupiter Trojans thus represent a rather primitive and isolated population. Several recent analyses confirmed the presence of several families among Trojans (e.g. Nesvorný et al. 2015, Vinogradova, 2015). The Trojan region as such is very favourable for dynamical studies of asteroid families, because there is no significant systematic Yarkovsky drift in semimajor axis due to the resonant dynamics. On the other hand, we have to be aware of boundaries of the libration zone, because ballistic transport can cause a partial depletion of family members. At the same time, as we have already shown in Bro\v{z} \& Rozehnal (2011), no family can survive either late phases of a slow migration of Jupiter, or Jupiter ``jump'', that results from relevant scenarios of the Nice model (Morbidelli et al. 2010). We thus focus on post-migration phase in this paper. We feel the need to evaluate again our previous conclusions on even larger datasets, that should also allow us to reveal as-of-yet unknown structures in the space of proper elements or unveil possible relations between orbital and physical properties (e.g. albedos, colours, diameters) of Jovian Trojans. In Section \ref{sec:obsdata} we use new observational data to compute appropriate resonant elements. In Section \ref{sec:physchar} we use albedos obtained by Grav et al. (2012) to derive size-frequency distributions and distribution of albedos, which seem to be slightly dependent on the proper inclination $I_{\rm p}$. In Section \ref{sec:groups} we identify families among Trojans with our new ``randombox'' method. We discuss properties of statistically significant families in Section \ref{sec:props}. Then we focus mainly on the Hektor family because of its unique D--type taxonomical classification, which is the first of its kind. We also discuss its long-term dynamical evolution. In Section \ref{sec:coll_model} we simulate collisional evolution of Trojans and we estimate the number of observable families among Trojans. Finally, in Section \ref{sec:SPH} we simulate an origin of the Hektor family using smoothed-particle hydrodynamics and we compare results for single and bilobed targets. Section \ref{sec:conclusions} is devoted to Conclusions.
\label{sec:conclusions} In this paper, we updated the list of Trojans and their proper elements, what allowed us to update parameters of Trojan families and to discover a new one (namely $2001\,\rm UV_{209}$ in $L_5$ population). We focused on the Hektor family, which seems the most interesting due to the bilobed shape of the largest remnant with a small moon and also its D-type taxonomical classification, which is unique among the collisional families observed so far. At the current stage of knowledge, it seems to us there are no major inconsistencies among the observed number of Trojan families and their dynamical and collisional evolution, at least in the current environment. As usual, we ``desperately'' need new observational data, namely in the size range from 5 to 10 km, which would enable us to constrain the ages of asteroid families on the basis of collisional modeling and to decide between two proposed ages of Hektor family, 1~to~4~Gyr or 0.1~to~2.5~Gyr. As expected, there are qualitative differences in impacts on single and bilobed targets. In our setup, the shockwave does not propagate easily into the secondary, so that only one half the mass is totally damaged as one can see in Figure \ref{fig:SPH}. On the other hand, the resulting SFDs are not that different, as we would expect. Even so, there is a large parameter space, which is still not investigated (i.e. the impact geometry with respect to the secondary, secondary impacts, the position in the orbit). SPH simulations of impacts on bilobed or binary targets thus seem very worthy for future research. Our work is also a strong motivation for research of disruptions of weak bodies (e.g. comets), better understanding the cometary disruption scaling law and also for experimental determination of material constants, which appear in the respective equation of state. As a curiosity, we can also think of searching for the remaining projectile, which could be still present among Trojans on a trajectory substantially different from that of family. A substantial part of projectile momentum is preserved in our simulations, so we may turn the logic and we may assume the projectile most likely came from the Trojan region and then it should remain in this region too.
16
7
1607.04677
In this work, we analyse Jovian Trojans in the space of suitable resonant elements and we identify clusters of possible collisional origin by two independent methods: the hierarchical clustering and a so-called randombox. Compared to our previous work, we study a twice larger sample. Apart from Eurybates, Ennomos and 1996 RJ families, we have found three more clusters - namely families around asteroids (20961) Arkesilaos, (624) Hektor in the L<SUB>4</SUB> libration zone and (247341) 2001 UV<SUB>209</SUB> in L<SUB>5</SUB>. The families fulfill our stringent criteria, I.e. a high statistical significance, an albedo homogeneity and a steeper size-frequency distribution than that of background. In order to understand their nature, we simulate their long term collisional evolution with the Boulder code and dynamical evolution using a modified SWIFT integrator. Within the framework of our evolutionary model, we were able to constrain the age of the Hektor family to be either 1-4 Gyr or, less likely, 0.1-2.5 Gyr, depending on initial impact geometry. Since (624) Hektor itself seems to be a bilobed-shape body with a satellite, I.e. an exceptional object, we address its association with the D-type family and we demonstrate that the moon and family could be created during a single impact event. We simulated the cratering event using a smoothed particle hydrodynamics. This is also the first case of a family associated with a D-type parent body.
false
[ "family", "initial impact geometry", "possible collisional origin", "dynamical evolution", "background", "suitable resonant elements", "clusters", "a smoothed particle hydrodynamics", "a single impact event", "SWIFT", "the Hektor family", "their long term collisional evolution", "a modified SWIFT integrator", "the D-type family", "L<SUB>5</SUB", "a D-type parent body", "1996 RJ families", "Boulder", "Jovian Trojans", "asteroids" ]
8.094302
16.276896
-1
12593934
[ "Tang, Xiaping", "Chevalier, Roger A." ]
2017MNRAS.465.3793T
[ "Shock evolution in non-radiative supernova remnants" ]
28
[ "Max Planck Institute for Astrophysics, Karl-Schwarzschild-Str 1, D-85741 Garching, Germany", "Department of Astronomy, University of Virginia, PO Box 400325, Charlottesville, VA 22904-4325, USA" ]
[ "2017A&A...597A.106A", "2017IAUS..331..206K", "2018ApJ...861..150P", "2018ApJ...866...40M", "2019ApJ...877..136F", "2019ApJ...887..198M", "2019MNRAS.488..978J", "2019NCimR..42..549B", "2019PASJ...71...61S", "2020A&A...644A..72P", "2020APh...12302492C", "2020ApJ...901..119R", "2021ApJ...907..111Z", "2021ApJ...907..117T", "2021MNRAS.505..755P", "2021MNRAS.505.4669J", "2021MNRAS.505.5301R", "2021MNRAS.508.3194B", "2021PhRvE.104d5213R", "2022ApJ...929...57V", "2022ApJ...930...28C", "2022MNRAS.511.3321C", "2022Sci...376...77H", "2023A&A...672A.194R", "2023JSP...190..118S", "2023MNRAS.523..132D", "2023MNRAS.523.1661N", "2023MNRAS.526.6214S" ]
[ "astronomy" ]
12
[ "shock waves", "methods: analytical", "ISM: supernova remnants", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1946RSPSA.186..273T", "1959sdmm.book.....S", "1963idp..book.....P", "1967dmiv.book.....R", "1982ApJ...258..790C", "1984ApJ...281..682H", "1985Ap&SS.112..225N", "1988ApJ...334..252C", "1988RvMP...60....1O", "1994ApJ...420..268C", "1994ShWav...4....1B", "1996sssi.book.....B", "1998JCoPh.144...70C", "1999ApJS..120..299T", "2000ApJS..128..403T", "2003ApJ...597..347L", "2012ApJ...746..130H", "2016A&A...590A..65M" ]
[ "10.1093/mnras/stw2978", "10.48550/arXiv.1607.06391" ]
1607
1607.06391_arXiv.txt
In this work, we derive simple analytical formulae characterizing the shock evolution in a non-radiative supernova remnant (SNR). In order to obtain an analytical solution, we constrain our discussion to a simple situation: spherical expansion in a smooth medium (no clouds) with negligible external thermal pressure. Thermal conduction, magnetic fields and acceleration of cosmic ray particles are also neglected for simplicity. The possible extension of the current model to more complicated situations will be studied in future work. The work presented here focuses on the evolution of a remnant in the post supernova phase. A point explosion with an ejecta mass of $M_{ej}$ and total energy of $E_{SN}$ is assumed as our initial condition. The energetic ejecta released in the supernova explosion drive a blast wave into the surrounding ambient medium, which is assumed to have a power law density profile with index $s$, i.e. $\rho_a \propto R^{-s}$. During the interaction between ejecta and ambient medium, both a forward shock into the surrounding medium with radius $R_b$ and a reverse shock into the expanding ejecta with radius $R_r$ are generated. The interface between the ejecta and the ambient medium is the shock contact discontinuity (CD). Its radius is defined as $R_c$. At early times, when the ejecta mass $M_{ej}\gg M_{sw}$, the swept up mass, the evolution of the remnant can follow two different evolutionary tracks depending on the spatial density distribution of the ejecta envelope. If the ejecta envelope has a shallow density profile $\rho \propto R^{-n}$ with power law index $n<5$, the early evolution of a SNR is characterized by the free expansion (FE) of the ejecta, with a narrow outer shocked region. In the FE solution, the CD expands freely with a constant velocity while the forward shock follows a similarity solution \citep{Parker63,HS84}. Due to the accumulation of shocked ambient medium ahead of the CD, it is found that $R_b=q_{b}R_c$ where $q_{b}>1$ is a dimensionless constant. If the ejecta envelope has a steep density profile $\rho \propto R^{-n}$ with power law index $n>5$, the early evolution of a SNR is instead described by the self similar driven wave (SSDW) solution \citep{Chevalier82,nadezhin85}. In the SSDW solution, $R_c\propto t^{(n-3)/(n-s)}$ based on dimensional analysis while $R_b=q_{b}R_c$ and $R_r=q_{r}R_c$ where $q_{b}$ and $q_{r}$ are dimensionless constants. As the blast wave expands, $M_{sw}$ gradually increases with time and eventually becomes dynamically important. When $M_{sw} \gg M_{ej}$, the expanding SNR has already lost the memory of ejecta mass $M_{ej}$ and starts to follow the self similar Sedov-Taylor (ST) solution \citep{Taylor46,Sedov59} in which $R_b\propto t^{2/(5-s)}$. According to the asymptotic behaviors described above at $M_{ej}\gg M_{sw}$ and $M_{ej}\ll M_{sw}$, \citep[][hereafter TM99]{TM99} derive analytic approximations for the evolution of the forward shock and reverse shock in a non-radiative SNR with further dynamical considerations. In TM99, the solution for the forward shock contains two parts: {\it a general ED solution} for the ED phase and {\it a general ST solution} for the ST phase. The transition time $t_{ST}$ defined in TM99, which separates the {\it general ED solution} from the {\it general ST solution}, is slightly different from the time when $M_{ej}=M_{sw}$ and in many cases is obtained through fitting numerical simulations. The {\it general ED solution} asymptotically approaches the FE solution when $n<5$ and the SSDW solution when $n>5$ as $t\rightarrow 0$, and is extended to finite $t$ by assuming the pressure behind the blast wave is proportional to that behind the reverse shock. It has two different forms depending on whether the reverse shock is in the envelope or the core of the ejecta. The {\it general ST solution} approaches the ST solution as $t\rightarrow \infty$ and equals to the value of the {\it general ED solution} at $t_{ST}$. It has the form of an offset power law and is designed to smoothly connect the {\it general ED solution} and the ST solution. The solution for the reverse shock in TM99 also contains two parts. In the ED phase, the reverse shock position is derived by assuming the reverse shock radius is proportional to the forward shock radius. In the ST phase, the reverse shock is described by a solution with constant acceleration in the unshocked ejecta frame. TM99 applied their method to a power law density ambient medium with a focus on the uniform medium and a brief discussion about the wind density profile. \cite{L&H03}, \cite{HL12} and \cite{Micelotta16} then studied the wind density profile in more detail with the method described in TM99. However, their solutions are not compared to numerical simulations. \cite{L&H03} and \cite{HL12} also presented analytical approximate formulae for the fitting coefficients used in the TM99 solution. Since all these analytical solutions \citep{TM99,L&H03,HL12,Micelotta16} are based on the method of TM99, from here on we refer to the TM model as the combination of the above solutions. Here we present a new analytical method to derive approximate solutions describing the shock evolution in a SNR from the ED phase to the ST phase. The method is based on dimensional analysis and depends on only the asymptotic behaviors of the remnant, i.e. the FE solution and the SSDW solution for $M_{sw}\ll M_{ej}$ and the ST solution for $M_{sw}\gg M_{ej}$. Because no further assumptions about the dynamical structure of the remnant are required as in the TM model, the analytical approximations discussed here are much simpler than the TM model solutions. The method presented here could potentially be extended to other problems involving the transition between two adjacent asymptotic limits. In Section \ref{sec:method}, we develop the analytical approach used to derive the approximate solutions. Then we use the new method to study the evolution of the forward shock and CD in a non-radiative SNR. Analytical approximations for both ejecta envelope with a shallow density profile $n<5$ and a steep density profile $n>5$ are investigated in detail. In Section \ref{sec:comparison}, we summarize the analytical approximations for both forward shock and CD, and then compare them with numerical simulations. For the forward shock, we also compare our new solutions with those from the TM model. We focus on two particularly interesting cases: SNR evolution in the interstellar medium with a constant density profile and SNR evolution in circumstellar material with a wind density profile. A reader who is only interested in the final expressions of the analytical approximations can go directly to this section. In Section \ref{sec:RS}, application of our new method to the reverse shock is discussed briefly. A final discussion and summary are in Section \ref{sec:DS}.
{\label{sec:DS}} In fitting the CD, we assume that the CD asymptotically approaches the power law relation $R_c^* \sim ct^{*b}$. For a wind density profile ($s=2$), the values of $b$ and $c$ (see Table \ref{CDs2}) we found are almost constant for ejecta with different density profiles, which implies a universal asymptotic limit for the CD like the ST solution for the forward shock. For a uniform medium, the derived $b$ and $c$ show larger variations because the reflected wave driven by the reverse shock complicates the situation. Thus we have to apply an arbitrary upper cutoff $t_{lim}^*$ during the fitting, which could affect the values of $b$ and $c$. In this paper, we present a new approach to derive analytical approximations describing the shock evolution in a non-radiative SNR. The new approach depends on only the asymptotic behaviors of the remnant during its evolution and thus is greatly simplified compared with the TM model. We then use the new method to closely investigate the shock evolution in a non-radiative SNR in both the interstellar medium with a constant density profile and a circumstellar medium with a wind density profile. We focus on the study of the forward shock and CD while application to the reverse shock is also briefly discussed. We compare our new analytical approximation with numerical simulations and find that a few percent accuracy is achieved for all investigated cases. For the forward shock, we also compare our new solutions to the TM model. In a uniform ambient medium, our solutions are comparable to the TM model while for a wind density profile medium our solutions perform better, especially when the ejecta envelope has a steep density profile. In order to obtain the analytical solution, we made several simplifying assumptions. Possible extensions of the current solutions to more complicated situations will be studied in future work. The transition from the ST phase to the radiative phase in SNR evolution has been discussed in \cite{Cioffi88}. In the future, we would like to use the method developed here to investigate the problem.
16
7
1607.06391
We present a new analytical approach to derive approximate solutions describing the shock evolution in non-radiative supernova remnants (SNRs). We focus on the study of the forward shock and contact discontinuity while application to the reverse shock is only discussed briefly. The spherical shock evolution of an SNR in both the interstellar medium with a constant density profile and a circumstellar medium with a wind density profile is investigated. We compared our new analytical solution with numerical simulations and found that a few per cent accuracy is achieved. For the evolution of the forward shock, we also compared our new solution to previous analytical models. In a uniform ambient medium, the accuracy of our analytical approximation is comparable to that in Truelove &amp; McKee. In a wind density profile medium, our solution performs better than that in Micelotta, Dwek &amp; Slavin, especially when the ejecta envelope has a steep density profile. The new solution is significantly simplified compared to previous analytical models, as it only depends on the asymptotic behaviours of the remnant during its evolution.
false
[ "non-radiative supernova remnants", "a wind density profile medium", "previous analytical models", "approximate solutions", "a steep density profile", "a wind density profile", "a constant density profile", "amp", "SNRs", "McKee", "The spherical shock evolution", "our new analytical solution", "the shock evolution", "Slavin", "numerical simulations", "a circumstellar medium", "a uniform ambient medium", "Dwek", "Truelove", "Micelotta" ]
3.996166
4.553815
-1
12436771
[ "Trani, Alessandro Alberto", "Mapelli, Michela", "Spera, Mario", "Bressan, Alessandro" ]
2016arXiv160707438T
[ "Dynamics of tidally captured planets in the Galactic Center" ]
0
[ "-", "-", "-", "-" ]
null
[ "astronomy" ]
6
[ "Astrophysics - Astrophysics of Galaxies" ]
[ "1999MNRAS.310..745M", "2012NatCo...3.1049M", "2012Natur.481...51G", "2013ApJ...764..154D", "2014ApJ...783..131Y", "2015ApJ...801L..26Y", "2015ApJ...806..197M" ]
[ "10.48550/arXiv.1607.07438" ]
1607
1607.07438_arXiv.txt
Recent radio continuum observations suggest the presence of photoevaporating protoplanetary disks in the innermost $\sim{}0.1$ pc of the Galactic Center (GC, \citealt{yus15}), very close to the super-massive black hole (SMBH). % \citet{map15} showed that rogue planets are too faint to be observed in the GC with current facilities, unless their near-infrared luminosity is enhanced by photoevaporation combined with tidal disruption induced by the SMBH. A protoplanetary origin has been suggested for the dusty object G2, which has been observed orbiting the SMBH on an highly eccentric orbit ($e \sim 0.98$) with very small periapsis ($a \sim 200 \AU$, \citealt{gil12}): \citet{mur12} proposed that G2 is a low-mass star with a proto-planetary disk, while \citet{map15} suggested that the properties of G2 are consistent with a planetary embryo tidally captured by the SMBH. In this proceeding, we study the possibility that planets and protoplanets are tidally captured by the SMBH.
Figure~\ref{fig:gorb} shows the trajectory of one of the simulated planets. The planet completes several orbits around the star, even if its motion is strongly perturbed by the tidal field of the SMBH. As the star approaches its periapsis, the planet is tidally captured by the the SMBH. Figure \ref{fig:gorb2} shows the cumulative probability map of finding an unbound planet in the semi-major axis -- eccentricity phase space. No planet can match the orbits of the G2 cloud. In particular, none of the simulated planets can achieve a highly-eccentric orbit. In fact, the closest periapsis passage of an unbound planet in our simulations is $1750\rm\, AU$, a factor of $\sim 9$ larger than the periapsis passage of the G2 cloud. We speculate that perturbations from other stars in the disk may bring planets into nearly-radial orbits. In forthcoming studies we will investigate the effect angular momentum transport and scatterings with other stars on the dynamics of planets in the disk. \begin{figure}[t!] \centering \resizebox{1.\hsize}{!}{\includegraphics[clip=true]{gorb.eps}} % \caption{\footnotesize Cumulative probability map of semi-major axis and eccentricity of the simulated planets. Red cross: G2 cloud. Cyan dotted line: inner edge of the CW disk.} \label{fig:gorb2} \end{figure}
16
7
1607.07438
Recent observations suggest ongoing planet formation in the innermost parsec of our Galaxy. The super-massive black hole (SMBH) might strip planets or planetary embryos from their parent star, bringing them close enough to be tidally disrupted. We investigate the chance of planet tidal captures by running three-body encounters of SMBH-star-planet systems with a high-accuracy regularized code. We show that tidally captured planets have orbits close to those of their parent star. We conclude that the final periapsis distance of the captured planet from the SMBH will be much larger than 200 AU, unless its parent star was already on a highly eccentric orbit.
false
[ "ongoing planet formation", "planet tidal captures", "planets", "its parent star", "their parent star", "Galaxy", "planetary embryos", "SMBH-star-planet systems", "SMBH", "the captured planet", "tidally captured planets", "a high-accuracy regularized code", "The super-massive black hole", "Recent observations", "a highly eccentric orbit", "the innermost parsec", "the final periapsis distance", "our Galaxy", "three-body encounters", "200 AU" ]
8.252548
13.95943
-1
12509708
[ "Asano, Katsuaki", "Mészáros, Peter" ]
2016PhRvD..94b3005A
[ "Ultrahigh-energy cosmic ray production by turbulence in gamma-ray burst jets and cosmogenic neutrinos" ]
23
[ "Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8582, Japan", "Department of Astronomy &amp; Astrophysics; Department of Physics; Center for Particle &amp; Gravitational Astrophysics; Pennsylvania State University, University Park, Pennsylvania 16802, USA" ]
[ "2017ApJ...846L..28X", "2017PTEP.2017lA106H", "2017SSRv..207...63W", "2017arXiv171101163Z", "2018A&A...616A..57R", "2018ApJ...853...43X", "2018ApJ...861...31A", "2018ApJ...866...51K", "2018NPPP..297..217M", "2018SSRv..214...41B", "2018jwpw.book...63W", "2019Ap&SS.364..105H", "2019ApJ...877...71T", "2019PhLB..789..393Z", "2019supe.book..419B", "2020JPhCS1468a2090A", "2020PhRvD.102b3003D", "2021PhRvD.104f3020L", "2021PhRvD.104j3005Z", "2022PhRvD.106b3028B", "2023MNRAS.524.4950K", "2023SciA....9J2778.", "2024PhRvD.109f3006L" ]
[ "astronomy", "physics" ]
7
[ "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1968Ap&SS...2..171M", "1975MNRAS.172..557S", "1984A&A...136..227S", "1987PhR...154....1B", "1988SvAL...14..255P", "1993A&A...272..161R", "1995ApJ...446..699P", "1995ApJ...453..332J", "1995ApJ...453..883V", "1995ApJ...454...60N", "1995PhRvL..75..386W", "1996ApJ...456..422K", "1996ApJ...461L..37B", "1996ApJ...472..245J", "1999APh....10...47B", "1999ApJ...517..700L", "2000A&A...360..789S", "2000ApJ...530..292M", "2000PhRvL..85.1362B", "2001ApJ...556..479D", "2001ApJ...556L..37M", "2002MNRAS.337.1349R", "2003ApJ...586..356Z", "2004A&A...413..807K", "2004APh....20..429G", "2004ApJ...610..550P", "2004RvMP...76.1143P", "2005ApJ...628L...9I", "2005MNRAS.364L..42F", "2006A&A...453...47K", "2006A&A...457..763G", "2006ApJ...638..811C", "2006ApJ...642..995P", "2006ApJ...647..539B", "2006ApJ...651..142H", "2006PhRvL..97e1101M", "2006RPPh...69.2259M", "2007ApJ...669..684B", "2007ApJ...670L..77I", "2007ApJ...671.1726S", "2007PhRvD..76l3001M", "2007PhRvE..75a6311R", "2008ApJ...681.1725S", "2008PhRvD..78b3005M", "2009APh....30..306T", "2009ApJ...699..953A", "2009ApJ...699.1261M", "2009ApJ...705.1714A", "2010ApJ...720..742K", "2010ApJ...725.1137L", "2010JCAP...10..013K", "2010MNRAS.406.1944W", "2010MNRAS.406.2113C", "2010MNRAS.407.1033B", "2011ApJ...734...77I", "2012APh....35..767T", "2012ApJ...746..164M", "2012ApJ...752...29H", "2012JCAP...11..058G", "2012MNRAS.420..483G", "2012Natur.484..351A", "2012PhRvD..85d9901G", "2012PhRvL.108m5003H", "2013ApJ...768L...1A", "2013ApJ...772L...1M", "2013ApJ...772L...4G", "2013ApJ...774...41K", "2013ApJ...775...87D", "2013JCAP...09..008A", "2013JCAP...10..013M", "2013PhRvD..88l1301M", "2013PhRvL.110l1101Z", "2013PhRvL.110x1101B", "2013PhRvL.111m1102M", "2014ApJ...780...64A", "2014ApJ...785...54A", "2014ApJ...790L..21A", "2014ApJ...791...71L", "2014JCAP...09..043T", "2014MNRAS.442.3026P", "2014MNRAS.445..570P", "2014PhRvL.113o5005G", "2015APh....62...66B", "2015ApJ...805L...5A", "2015ApJ...806..167G", "2015ApJ...807..148G", "2015ApJ...808L..18A", "2015ApJ...809...98A", "2015MNRAS.449..551K", "2015MNRAS.454.2242A", "2015NatCo...6.6783B", "2016ApJ...825..122H", "2016ApJ...831L...8G", "2016EPJST.225.1071V", "2016MNRAS.460.4135P", "2017ApJ...836...47B", "2017ApJ...837...33B" ]
[ "10.1103/PhysRevD.94.023005", "10.48550/arXiv.1607.00732" ]
1607
1607.00732_arXiv.txt
Introduction} The origin of ultrahigh-energy cosmic rays (UHECRs) above the ankle energy ($\sim 10^{18.5}$ eV) is a matter of ongoing discussions. Although jets in active galactic nuclei (AGN) are the most widely considered candidates for the UHECR sources \citep{rac93}, the observed degree of anisotropy in the arrival distribution indicates a source density larger than $10^{-4}~\mbox{Mpc}^{-3}$ for a pure proton compositon ($10^{-6}~\mbox{Mpc}^{-3}$ for a pure iron composition) \citep{tak09,tak12}, which disfavors Fanaroff-Reily II galaxies and BL Lac objects. Other types of relatively low-luminosity AGNs like Seyfert galaxies may not satisfy the luminosity requirement ($\gtrsim 10^{46}~\mbox{erg}~\mbox{s}^{-1}$ for protons) needed for UHECR acceleration \citep{nor95}. Clusters of galaxies with strong accretion shocks are also candidates of UHECR sources \citep{kan96,ino05}. However, while the Telescope Array experiment reported a cluster of UHECR events \citep{abb14} in a $20^\circ$ radius (TA hot spot), there is no clear excess in the direction toward the nearby massive cluster, Virgo. This situation motivates us to revisit gamma-ray bursts (GRBs) as UHECR sources \citep{vie95,wax95}, although the GRB hypothesis has been considered to have disadvantages. The prompt emission of GRBs is believed to be emitted from collimated ultrarelativistic outflows. In most of the GRB UHECR models, the internal shocks formed in the GRB outflows \citep{pir04,mes06} are supposed to be the UHECR acceleration site. In this case, the shock accelerated particles (hereafter we assume UHECR to be protons) form a power-law number spectrum ($N(\varepsilon) \propto \varepsilon^{-p}$) with a typical index of $p \sim 2$. The observed GRB rate \footnote{A narrow (wide) average jet opening angle leads to a high (low) actual GRB rate but low (high) energy release per burst. Thus, roughly speaking, the actual energy release rate from GRBs is independent of the jet opening angle. Hereafter, we adopt the observed GRB rate with the isotropically equivalent energy release.} is so low that the required cosmic-ray luminosity to agree with the observed UHECR flux is 30--100 times the gamma-ray luminosity or more \citep[see e.g.,][]{mur08,bae15}. Such a high luminosity seems unfavorable in the light of the available energy budget. Assuming a cosmic-ray luminosity much larger than the gamma-ray luminosity, the secondary neutrino flux has been calculated by many authors \citep[e.g., Refs.][]{gue04,mur08,bae15,bus15,bus16}. However, the IceCube neutrino telescope has detected no significant high-energy neutrino emission associated with classical GRBs \citep{abb12,gao13,aar15-GRB}. This severely constrains the UHECR luminosity in GRBs, although a moderate value of the ratio of the UHECR to gamma-ray luminosity ($f_{\rm CR}\sim10$) is allowed \citep{He12,asa14}. In particular, a GRB UHECR model with the moderate ratio $f_{\rm CR}=10$ can reproduce the observed UHECR flux, but only above $\sim 10^{20}$ eV \citep{asa14}. Furthermore, while most of the previous studies of the GRB neutrino emission associated with UHECR have omitted a discussion of the secondary gamma rays, the required high UHECR luminosity must result in a spectral shape of the gamma rays which differs from the typically observed ones, as result of the hadronic cascade initiated by the collisions with gamma-ray photons \citep{asa09,asa14,pet14}. In this paper, we discuss a different scenario for the UHECR production in GRBs, which may avoid the above-mentioned problems. Among the latter, the major difficulty arises because the power-law spectrum of index $p=2$ with an exponential cutoff, which has been frequently assumed in the shock acceleration model, leads to a large energy fraction residing below the ankle energy. On the other hand, if the spectral index were shallower than 2, then most of the cosmic-ray energy would be concentrated around the highest energy range. This could reduce the total proton energy budget, putting the bulk of the UHECR energy above the ankle. Such a hard spectrum would probably involve a different acceleration mechanism or acceleration site from those in the internal shock models. Another requirement in the GRB UHECR model is the suppression of the secondary gamma rays and neutrinos, the first being constrained by Fermi and the second by IceCube observations. This suggests that the UHECR acceleration site should be significantly outside the usual photon emission site, in order to reduce the photopion production efficiency. Such a setup could provide a convincing solution for avoiding an overly luminous gamma-ray/neutrino emission compatible with the required large UHECR luminosity, unless the average bulk Lorentz factor of the GRB jets is $\gtrsim 1000$ \citep{asa09}, or the gamma-ray emission site is also located at a larger distance than usually assumed \citep{pet14b,bus16}. In some observational \citep{gui15,gui16} and theoretical \citep[e.g., Ref.][]{bus16} studies, it has been argued that multiple spectral components of the prompt $\gamma$-ray emission may arise from different emission sites or mechanisms. This encourages us to consider a different site for the UHECR acceleration. An alternative model for the prompt gamma-ray emission is the dissipative photosphere model, which has been discussed by Refs. \citep{mes00,gia11,pee06,iok07,laz10,bel10,asa13}. In this model, the photon emission site is at a distance $\sim 10^{10}-10^{13}$ cm from the central engine, leaving a large fraction of the bulk kinetic energy of the flow to be dissipated at a larger distance. Numerical simulations of the deceleration of such relativistic outflows \citep{duf13} show the development of a Rayleigh-Taylor instability at large distances $\gtrsim 10^{16}$ cm. In such regions, stochastic acceleration via turbulence can accelerate UHECRs, and the photopion production efficiency will be significantly low. The stochastic acceleration can yield a hard UHECR spectrum with $p<2$ \citep{sch84,par95,bec06,sta08}, having been discussed as a possible electron acceleration mechanism in AGN jets \citep{bot99,sch00,kat06,kak15}. Stochastic acceleration of electrons via turbulence has also been discussed in connection with the mechanism of the GRB prompt gamma-ray emission \citep{Byk96,der01,asa09b,mur12,asa15b}. In particular, recent numerical simulations of the stochastic acceleration and photon emission in AGN jets \citep{asa14b,asa15} succeed in reproducing the wide-band spectra from radio to gamma rays and the gamma-ray light curves, showing that stochastic acceleration in relativistic outflows is an attractive option. In this paper, we propose a UHECR production model in GRB outflows based on the stochastic acceleration of protons via turbulence at a large distance, well outside the photon emission site. As will be shown, this optimized model can avoid both the problems of an overly large UHECR loading and an overly luminous secondary gamma-ray/neutrino emission.
We propose a possible model of the UHECR production by the stochastic proton acceleration via turbulence in the GRB jets. The UHECR spectrum at injection is harder than in previous models and shows a curved feature, which does not conflict with the observed UHECR spectral shape, its presence being felt mainly above the ankle. The required typical cosmic-ray luminosity is $\sim 10^{53.5}~ \mbox{erg}~\mbox{s}^{-1}$, which is moderate compared to previous GRB UHECR models. An overly luminous secondary gamma-ray/neutrino emission initiated by photopion production is avoided because the acceleration site is expected to be well outside the photon emission radius. A predicted hard spectrum of GZK neutrinos in the $10^{17}$--$10^{18}$ eV range can be a clue to constraining the parent UHECR spectrum. The UHECR spectrum at injection is very sensitive to the model parameters, which are uncertain and may have a substantial dispersion. Especially the $L_\gamma$-$f_{\rm CR}$ or $L_\gamma$-$\Gamma$ relations are not well defined. Depending on those parameters, the dominant UHECR contribution may come from the relatively low-luminosity GRBs ($L_\gamma<10^{52.5}~\mbox{erg}~\mbox{s}^{-1}$) or vice versa. Although we cannot, so far, predict a quantitatively precise UHECR spectrum, the possibility of a hard spectrum such as discussed in this paper appears to be an attractive idea for overcoming the difficulties in the GRB UHECR hypothesis.
16
7
1607.00732
We propose a novel model to produce ultrahigh-energy cosmic rays (UHECRs) in gamma-ray burst jets. After the prompt gamma-ray emission, hydrodynamical turbulence is excited in the GRB jets at or before the afterglow phase. The mildly relativistic turbulence stochastically accelerates protons. The acceleration rate is much slower than the usual first-order shock acceleration rate, but in this case it can be energy independent. The resultant UHECR spectrum is so hard that the bulk energy is concentrated in the highest energy range, resulting in a moderate requirement for the typical cosmic-ray luminosity of ∼1 0<SUP>53.5</SUP> erg s<SUP>-1</SUP> . In this model, the secondary gamma-ray and neutrino emissions initiated by photopion production are significantly suppressed. Although the UHECR spectrum at injection shows a curved feature, this does not conflict with the observed UHECR spectral shape. The cosmogenic neutrino spectrum in the 10<SUP>17</SUP>- 10<SUP>18</SUP> eV range becomes distinctively hard in this model, which may be verified by future observations.
false
[ "future observations", "ultrahigh-energy cosmic rays", "gamma-ray burst jets", "∼1", "photopion production", "hydrodynamical turbulence", "the highest energy range", "the observed UHECR spectral shape", "UHECR", "the afterglow phase", "the typical cosmic-ray luminosity", "protons", "UHECRs", "GRB", "the usual first-order shock acceleration rate", "the bulk energy", "the prompt gamma-ray emission", "first", "the GRB jets", "the secondary gamma-ray and neutrino emissions" ]
5.727521
0.270598
13
3817758
[ "Guzik, J. A.", "Houdek, G.", "Chaplin, W. J.", "Smalley, B.", "Kurtz, D. W.", "Gilliland, R. L.", "Mullally, F.", "Rowe, J. F.", "Bryson, S. T.", "Still, M. D.", "Antoci, V.", "Appourchaux, T.", "Basu, S.", "Bedding, T. R.", "Benomar, O.", "Garcia, R. A.", "Huber, D.", "Kjeldsen, H.", "Latham, D. W.", "Metcalfe, T. S.", "Pápics, P. I.", "White, T. R.", "Aerts, C.", "Ballot, J.", "Boyajian, T. S.", "Briquet, M.", "Bruntt, H.", "Buchhave, L. A.", "Campante, T. L.", "Catanzaro, G.", "Christensen-Dalsgaard, J.", "Davies, G. R.", "Doğan, G.", "Dragomir, D.", "Doyle, A. P.", "Elsworth, Y.", "Frasca, A.", "Gaulme, P.", "Gruberbauer, M.", "Handberg, R.", "Hekker, S.", "Karoff, C.", "Lehmann, H.", "Mathias, P.", "Mathur, S.", "Miglio, A.", "Molenda-Żakowicz, J.", "Mosser, B.", "Murphy, S. J.", "Régulo, C.", "Ripepi, V.", "Salabert, D.", "Sousa, S. G.", "Stello, D.", "Uytterhoeven, K." ]
2016ApJ...831...17G
[ "Detection of Solar-like Oscillations, Observational Constraints, and Stellar Models for θ Cyg, the Brightest Star Observed By the Kepler Mission" ]
15
[ "Los Alamos National Laboratory, XTD-NTA, MS T-082, Los Alamos, NM 87545, USA", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; School of Physics and Astronomy, University of Birmingham, Birmingham B15 2TT, UK", "Astrophysics Group, School of Physical and Geographical Sciences, Lennard-Jones Laboratories, Keele University, Staffordshire, ST5 5BG, UK", "Jeremiah Horrocks Institute, University of Central Lancashire, Preston PR1 2HE, UK", "Center for Exoplanets and Habitable Worlds, The Pennsylvania State University, University Park, PA 16802, USA", "SETI Institute/NASA Ames Research Center, Moffett Field, CA 94035, USA", "SETI Institute/NASA Ames Research Center, Moffett Field, CA 94035, USA", "NASA Ames Research Center, Bldg. 244, MS-244-30, Moffett Field, CA 94035, USA", "NASA Ames Research Center, Bldg. 244, MS-244-30, Moffett Field, CA 94035, USA; Bay Area Environmental Research Institute, 560 Third Street W., Sonoma, CA 95476, USA", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark", "Institut d'Astrophysique Spatiale, Universitè de Paris Sud–CNRS, Batiment 121, F-91405 ORSAY Cedex, France", "Department of Astronomy, Yale University, P.O. Box 208101, New Haven, CT 06520-8101, USA", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Sydney Institute for Astronomy (SIfA), School of Physics, University of Sydney, NSW 2006, Australia", "Sydney Institute for Astronomy (SIfA), School of Physics, University of Sydney, NSW 2006, Australia ; NYUAD Institute, Center for Space Science, New York University Abu Dhabi, P.O. Box 129188, Abu Dhabi, UAE", "Laboratoire AIM, CEA/DRF—CNRS—Univ. Paris Diderot—IRFU/SAp, Centre de Saclay, F-91191 Gif-sur-Yvette Cedex, France", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Sydney Institute for Astronomy (SIfA), School of Physics, University of Sydney, NSW 2006, Australia", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO 80301, USA", "Instituut voor Sterrenkunde, KU Leuven, Celestijnenlaan 200D, B-3001 Leuven, Belgium", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Sydney Institute for Astronomy (SIfA), School of Physics, University of Sydney, NSW 2006, Australia ; Australian Astronomical Observatory, P.O. Box 915, North Ryde, NSW 1670, Australia", "Instituut voor Sterrenkunde, KU Leuven, Celestijnenlaan 200D, B-3001 Leuven, Belgium; Department of Astrophysics/IMAPP, Radboud University Nijmegen, 6500 GL Nijmegen, The Netherlands", "Université de Toulouse, UPS-OMP, IRAP, F-65000, Tarbes, France", "Department of Astronomy, Yale University, P.O. Box 208101, New Haven, CT 06520-8101, USA", "Institut d'Astrophysique et de Géophysique, Université de Liège, Quartier Agora, Allée du 6 août 19C, B-4000, Liège, Belgium", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Aarhus Katedralskole, Skolegyde 1, DK-8000 Aarhus C, Denmark", "Niels Bohr Institute, University of Copenhagen, DK-2100 Copenhagen, Denmark ; Centre for Star and Planet Formation, Natural History Museum of Denmark, University of Copenhagen, DK-1350 Copenhagen, Denmark", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; School of Physics and Astronomy, University of Birmingham, Birmingham B15 2TT, UK", "INAF-Osservatorio Astrofisico di Catania, Via S.Sofia 78, I-95123 Catania, Italy", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; School of Physics and Astronomy, University of Birmingham, Birmingham B15 2TT, UK ; Laboratoire AIM, CEA/DRF—CNRS—Univ. Paris Diderot—IRFU/SAp, Centre de Saclay, F-91191 Gif-sur-Yvette Cedex, France", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Department of Geoscience, Aarhus University, Hoegh-Guldbergs Gade 2, DK-8000, Aarhus C, Denmark ; High Altitude Observatory, National Center for Atmospheric Research, P.O. Box 3000, Boulder, CO 80307, USA", "The Department of Astronomy and Astrophysics, University of Chicago, 5640 S. Ellis Avenue, Chicago, IL 60637, USA", "Astrophysics Group, School of Physical and Geographical Sciences, Lennard-Jones Laboratories, Keele University, Staffordshire, ST5 5BG, UK; Department of Physics, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; School of Physics and Astronomy, University of Birmingham, Birmingham B15 2TT, UK", "INAF-Osservatorio Astrofisico di Catania, Via S.Sofia 78, I-95123 Catania, Italy", "Apache Point Observatory, Sloan Digital Sky Survey, P.O. Box 59, Sunspot, NM 88349, USA ; New Mexico State University, Department of Astronomy, P.O. Box 30001, Las Cruces, NM 88003-4500, USA", "Institute for Computational Astrophysics, Department of Astronomy and Physics, Saint Mary's University, Halifax, NS B3H 3C3, Canada", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Max Planck Institute for Solar System Research, SAGE Research Group, Justus-von-Liebig-Weg 3, D-37077 Gttingen, Germany", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Department of Geoscience, Aarhus University, Hoegh-Guldbergs Gade 2, DK-8000, Aarhus C, Denmark", "Thüringer Landessternwarte Tautenburg (TLS), Sternwarte 5, D-07778 Tautenburg, Germany", "Université de Toulouse, UPS-OMP, IRAP, F-65000, Tarbes, France; CNRS, IRAP, 57 Avenue d'Azereix, BP 826, F-65008, Tarbes, France", "Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO 80301, USA", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; School of Physics and Astronomy, University of Birmingham, Birmingham B15 2TT, UK", "Instytut Astronomiczny Uniwersytetu Wrocławskiego, ul. Kopernika 11, 51-622 Wrocław, Poland", "LESIA—Observatoire de Paris/CNRS, Sorbonne Universités, UMC Univ. Paris 06, Univ. Paris Diderot, Sorbonne Paris Cité, France", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Sydney Institute for Astronomy (SIfA), School of Physics, University of Sydney, NSW 2006, Australia", "Instituto de Astrofísica de Canarias, E-38205, La Laguna, Tenerife, Spain ; Universidad de La Laguna, Dpto de Astrofísica, E-38206, Tenerife, Spain", "INAF-Osservatorio Astronomico di Capodimonte, Via Moiariello 16, I-80131 Napoli, Italy", "Laboratoire AIM, CEA/DRF—CNRS—Univ. Paris Diderot—IRFU/SAp, Centre de Saclay, F-91191 Gif-sur-Yvette Cedex, France", "Instituto de Astrofísica e Ciências do Espaço, Universidade do Porto, CAUP, Rua das Estrelas, 4150-762 Porto, Portugal", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark; Sydney Institute for Astronomy (SIfA), School of Physics, University of Sydney, NSW 2006, Australia", "Instituto de Astrofísica de Canarias, E-38205, La Laguna, Tenerife, Spain ; Universidad de La Laguna, Dpto de Astrofísica, E-38206, Tenerife, Spain" ]
[ "2017A&A...602A..32A", "2017MNRAS.471.2882W", "2018ARA&A..56...83B", "2018ApJ...865L..20F", "2019ApJS..244...18P", "2019ApJS..245....8P", "2019LRSP...16....4G", "2019MNRAS.485..560C", "2020ApJ...905..108T", "2020svos.conf..457H", "2021A&ARv..29....4S", "2021NewA...8401522K", "2022A&A...658A..47K", "2023MNRAS.523.1441R", "2024ApJS..271...55K" ]
[ "astronomy" ]
8
[ "asteroseismology", "stars: fundamental parameters", "stars: interiors", "stars: solar-type", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.3847/0004-637X/831/1/17", "10.48550/arXiv.1607.01035" ]
1607
1607.01035.txt
\label{sec:introduction} The mission of the NASA {\it Kepler} spacecraft, launched 2009 March 7, was to search for Earth-sized planets around Sun-like stars in a fixed field of view in the Cygnus-Lyra region using high-precision CCD photometry to detect planetary transits \citep{Borucki2010}. As a secondary mission {\it Kepler} surveyed and monitored over 10,000 stars for asteroseismology, using the intrinsic brightness variations caused by pulsations to infer the star's mass, age, and interior structure \citep{2010PASP..122..131G}. %\footnote{See http://kepler.nasa.gov for more information about the {\it Kepler} %mission.} After the failure of the second of four reaction wheels, the {\it Kepler} mission transitioned into a new phase, K2 \citep{2014PASP..126..398H}, observing fields near the ecliptic plane for about 90 days each, with a variety of science objectives including planet searches. The $V\,=\,4.48$ F3 spectral-type main-sequence star $\theta$\,Cyg, also known as 13\,Cyg, HR\,7469, HD\,185395, 2MASS 19362654+5013155, HIP 96441, and KIC\,11918630, where KIC = {\it Kepler} Input Catalog \citep{2011AJ....142..112B}, is the brightest star that fell on active pixels in the original {\it Kepler} field of view. $\theta$~Cyg is nearby and bright, so that high-precision ground-based data can be combined with high signal-to-noise and long time-series {\it Kepler} photometry to provide constraints for asteroseismology. The position of $\theta$~Cyg in the HR~diagram is near that of known $\gamma$\,Dor pulsators, suggesting the possibility that it may exhibit high-order gravity mode pulsations, which would probe the stellar interior just outside its convective core. $\theta$~Cyg is also cool enough to exhibit solar-like $p$-mode (acoustic) oscillations, which probe both the interior and envelope structure. \begin{figure}%[tb] \center{\includegraphics[width=\columnwidth]{Kepler_FOV.eps}} \caption{{\it Kepler} field of view with stars marked according to stellar magnitude created using ``The Sky'' astronomy software (http://www.bisque.com/sc/pages/TheSkyX-Editions.aspx). The location of $\theta$~Cyg is shown by the red arrow, and is marked with the symbol $\theta$ and a filled circle designating magnitude 4-5. Note that all stars brighter than $\theta$ Cyg fall in the regions between the CCD arrays to avoid saturating pixels.} \label{KeplerFOV} \end{figure} %The revised Hipparcos parallax of $54.54 \pm 0.15$ milliarcseconds \citep{vanLeeuwen2007b} gives a distance of $18.33\,\pm\ %0.05$\,pc, which, combined with an estimate of its apparent bolometric luminosity \citep{2008ApJS..176..276V}, gives L = 0.63\,\pm\, %0.03$\,L$_{\odot}$. Other parameters in the literature derived from multi-color photometry and %high-resolution spectroscopy are $T_{\rm eff} = 6745\,\pm\,150 $\,K, $\log g = 4.2 \pm %50.2$ (cgs units), and $\rm{[M/H]} = -0.04$ \citep{Erspamer03, 2003AJ....126.2048G}. %5\citet{Nordstrom04} give a mass $1.38\,\pm\,0.05$\,M$_{\odot}$, and age 1.5 %+0.6/-0.7\,Gyr using Str\"omgren photometry plus the Padova stellar model grid. %Appendix A summarizes results of ground-based observations from the literature and the field around $\theta$ Cyg. Section \ref{sec:spectroscopy} describes analyses of new spectra of $\theta$ Cyg. $\theta$ Cyg has been observed using adaptive optics \citep{2009A&A...506.1469D}. It has a resolved binary M-dwarf companion of $\sim$0.35\,M$_{\odot}$ with separation 46\,AU. Following the orbit for nearly an orbital period (unfortunately $\sim$~230 y) will eventually give an accurate dynamical mass for $\theta$ Cyg. Also, the system shows a 150-d quasi-period in radial velocity, suggesting that one or more planets could accompany the stars \citep{2009A&A...506.1469D}. $\theta$ Cyg has also been observed using optical interferometry \citep[][see Section \ref{sec:interferometry}]{ 2012A&A...545A...5L, 2012ApJ...746..101B, 2013MNRAS.tmp.1445W}. These observations provide tight constraints on the radius of $\theta$ Cyg and therefore a very useful constraint for asteroseismology. $\theta$ Cyg's projected rotational velocity is low; $v~\sin i$ = $3.4 \pm 0.4$ \,km s$^{-1}$ \citep[][see Section \ref{sec:spectroscopy}]{1984ApJ...281..719G}. If $\sin i$ is not too small, $\theta$ Cyg's slow rotation should simplify mode identification and pulsation modeling, as spherical approximations and low-order perturbation theory for the rotational splitting should be adequate. This paper is intended to provide background on the $\theta$~Cyg system and to be a first look at the {\it Kepler} photometry data and consequences for stellar models and asteroseismology. We present light curves and detection of the solar-like $p$-modes based on {\it Kepler} data taken in observing Quarters 6 and 8 (Section \ref{sec:detection}). We summarize ground-based observational constraints from the literature (Appendix A) and present analyses based on new spectroscopic observations (Section \ref{sec:spectroscopy}) and optical interferometry (Section \ref{sec:interferometry}). We discuss inference of stellar parameters based on the large separation and frequency of maximum amplitude (Section \ref{sec:stellarparameters}), line widths (Section \ref{sec:damping}), and mode identification (Section \ref{sec:ModeID}). We use the observed $p$-mode oscillation frequencies and mode identifications as constraints for stellar models using several methods (Section \ref{sec:models}). We discuss predictions for $\gamma$ Dor $g$-mode pulsations (Section \ref{sec:gmodepredictions}), and results of a search for low frequencies consistent with $g$ modes (Section \ref{sec:gmodesearch}). We conclude with motivation for continued study of $\theta$ Cyg (Section \ref{sec:conclusions}). We do not include in this paper the analyses of data from Quarters 12-17 for several reasons. First, we completed the bulk of this paper, including the spectroscopic analyses, and first asteroseismic analyses at the time when only the Q6 and Q8 data were available. Second, a problem has emerged with the {\it Kepler} data reduction pipeline for the latest data release for short-cadence data\footnote{https://archive.stsci.edu/kepler/KSCI-19080-002.pdf} that will not be corrected until later in 2016; while $\theta$ Cyg is not on the list of affected stars, because $\theta$ Cyg required so many pixels and special processing, more work is needed to confirm that the problem has not introduced additional noise in the light curve. We estimate that inclusion of the full time-series data will result in finding a few more frequencies, and will improve the precision of the frequencies obtained by a factor of $\sim$1.8. %(T. Appourchaux, private communication, 2016) Comparison of studies of the bright (V = 5.98) {\it Kepler} targets 16 Cyg A and B using one month versus thirty months of data show that the longer time series improved the accuracy and precision of results, but did not significantly change the frequencies or inferred stellar model parameters \citep{2012ApJ...748L..10M,2015ApJ...811L..37M}. Detailed analyses making use of the remaining time-series data and the {\it Kepler} pixel data will be the subject of future papers.
\label{sec:conclusions} We have analyzed Quarters 6 and 8 of {\it Kepler} $\theta$ Cyg data, finding solar-like $p$-modes, and not finding $\gamma$ Dor gravity modes that were initially expected given $\theta$ Cyg's spectral tye. We have obtained new ground-based spectroscopic and interferometric observations and updated the observational constraints. Stellar models of $\theta$ Cyg that fit the $p$-mode frequencies and spectroscopic and interferometric constraints on $R$, $T_\mathrm{eff}$, log $g$, and [M/H] are predicted to show $g$-mode pulsations driven by the convective-blocking mechanism, according to nonadiabatic pulsation models. However, analysis of the light curves did not reveal any $g$ modes. Reprocessed {\it Kepler} observations of $\theta$ Cyg for Quarters 12 through 17 including the pipeline corrections will be available in late 2016. We intend to examine the pixel-by-pixel data to remove the background binary if possible. As noted by \cite{2013A&A...556A..52T} in analysis of their sample of 69 $\gamma$ Dor stars, use of the pixel data eliminated many spurious low frequencies detected using the standard pre-processed light curves. Analyses of a longer time series may reduce noise due to granulation, and more definitively rule out the presence of $g$ modes or identify features in the light curve resulting from rotation and stellar activity. The {\it Kepler} observations of $\theta$ Cyg, in conjunction with studies of many other A-F stars observed by {\it Kepler} and CoRoT, will be key to understanding the puzzles of $\gamma$ Dor/$\delta$ Sct hybrids and pulsating variables that appear to lie outside of instability regions expected from theoretical models, and to test stellar model physics and possible alternative pulsation driving mechanisms. Attempting to find $g$ modes in $\theta$ Cyg and other mid-F spectral type stars is worthwhile, as $g$ modes are more sensitive to the stellar interior near the convective core boundary than are $p$ modes. Seismic measurements of convective core size and shape, and the structure of the overshooting region will help reduce uncertainties in stellar ages and understand the roles of penetrative overshooting and diffusive mixing. Progress has already been made in this area for {\it Kepler} slowly-pulsating B stars that are $g$ mode pulsators by, e.g., \cite{2015A&A...580A..27M}, who used the spacings of 19 consecutive $g$ modes in KIC 9526294 to distinguish between models using exponentially decaying vs. a step-function overshooting prescription, and diagnose the need for additional diffusive mixing. However, note that progress is also being made studying convective cores using $p$ modes in low-mass stars \citep[see, e.g.][]{2016arXiv160302332D}, as the molecular weight gradient outside the convective core introduces a discontinuity in sound-speed profile that is diagnosable with $p$ modes. The core size and mode frequencies are also affected by rotation that is likely to be more rapid in the core than in the envelope. \cite{2015ApJS..218...27V} discuss $g$-mode periods and spacings for a sample of 67 $\gamma$ Dor stars observed by {\it Kepler}, and find correlations between $v~\sin i$, $T_{\rm eff}$, period spacing values, and dominant periods. van Reeth et al. (2016, submitted), discuss a method for mode identification of high-order $g$ modes from the period spacing patterns for $\gamma$ Dor stars, allowing to deduce rotation frequency near the core. \cite{2015EPJWC.10101005B} discuss using period \'echelle diagrams for {\it Kepler} $\gamma$ Dor stars to measure period spacings and identify rotationally split multiplets with $l$ = 1 and $l$ = 2. \cite{2015MNRAS.454.1792K} study KIC 10080943, two hybrid $\delta$ Sct/$\gamma$ Dor stars in a non-eclipsing spectroscopic binary, and are able to use rotational splitting to estimate core rotation rates. Because $\theta$ Cyg is nearby and bright, and data can be obtained with excellent precision, it is also a worthwhile target for continued long time-series ground- or space-based photometric or spectroscopic observations. With an even longer time series of data (obtainable by a follow-on to the {\it Kepler} mission), there is the possibility to study rotational splitting and differential rotation, infer convection zone depth directly from oscillation frequency inversions, measure sin~$i$ directly from amplitude differences of rotationally split modes, and investigate possible magnetic activity cycles.
16
7
1607.01035
θ Cygni is an F3 spectral type magnitude V = 4.48 main-sequence star that was the brightest star observed by the original Kepler spacecraft mission. Short-cadence (58.8 s) photometric data using a custom aperture were first obtained during Quarter 6 (2010 June-September) and subsequently in Quarters 8 and 12-17. We present analyses of solar-like oscillations based on Q6 and Q8 data, identifying angular degree l = 0, 1, and 2 modes with frequencies of 1000-2700 μHz, a large frequency separation of 83.9 ± 0.4 μHz, and maximum oscillation amplitude at frequency ν <SUB>max</SUB> = 1829 ± 54 μHz. We also present analyses of new ground-based spectroscopic observations, which, combined with interferometric angular diameter measurements, give T <SUB>eff</SUB> = 6697 ± 78 K, radius 1.49 ± 0.03 R <SUB>⊙</SUB>, [Fe/H] = -0.02 ± 0.06 dex, and log g = 4.23 ± 0.03. We calculate stellar models matching these constraints using the Yale Rotating Evolution Code and the Asteroseismic Modeling Portal. The best-fit models have masses of 1.35-1.39 M <SUB>⊙</SUB> and ages of 1.0-1.6 Gyr. θ Cyg’s T <SUB>eff</SUB> and log g place it cooler than the red edge of the γ Doradus instability region established from pre-Kepler ground-based observations, but just at the red edge derived from pulsation modeling. The pulsation models show γ Dor gravity modes driven by the convective blocking mechanism, with frequencies of 1-3 cycles per day (11 to 33 μHz). However, gravity modes were not seen in Kepler data; one signal at 1.776 cycles per day (20.56 μHz) may be attributable to a faint, possibly background, binary.
false
[ "maximum oscillation amplitude", "interferometric angular diameter measurements", "Kepler data", "g = 4.23 ±", "SUB", "g place", "Q8 data", "pulsation modeling", "pre-Kepler ground-based observations", "gravity modes", "± 78 K", "Quarters", "eff</SUB", "day", "frequencies", "new ground-based spectroscopic observations", "frequency", "83.9 ±", "Kepler", "the original Kepler spacecraft mission" ]
7.315479
11.354347
81
12472875
[ "Tritsis, Aris", "Tassis, Konstantinos" ]
2016MNRAS.462.3602T
[ "Striations in molecular clouds: streamers or MHD waves?" ]
36
[ "Department of Physics and ITCP, University of Crete, PO Box 2208, GR-71003 Heraklion, Greece", "Department of Physics and ITCP, University of Crete, PO Box 2208, GR-71003 Heraklion, Greece; IESL, Foundation for Research and Technology-Hellas, PO Box 1527, GR-71110 Heraklion, Crete, Greece" ]
[ "2017ApJ...847..140C", "2017arXiv170301542B", "2018MNRAS.481.5275T", "2018Sci...360..635T", "2019ASSP...55..117T", "2019ApJ...873...38T", "2019ApJ...878..157X", "2019FrASS...6....5H", "2019FrASS...6....7K", "2019MNRAS.485.4686W", "2019PASA...36...29A", "2020A&A...644A..27B", "2020ApJ...891...84C", "2020MNRAS.492..668B", "2020MNRAS.498.1593B", "2020MNRAS.499.4785K", "2020SSRv..216...76B", "2021A&A...647A.186S", "2021A&A...649A..33Y", "2021A&A...655A..76S", "2021MNRAS.500.2831W", "2021MNRAS.504.4354B", "2022A&A...665A..77S", "2022A&A...668A..42C", "2022FrASS...9.0900B", "2022MNRAS.513.1275A", "2022MNRAS.514.3593T", "2022arXiv221115586Z", "2023A&A...670A..59G", "2023A&A...673A..76S", "2023ASPC..534..153H", "2023ASPC..534..233P", "2023MNRAS.524.3201B", "2023MNRAS.526.4345K", "2023arXiv231020572M", "2024AJ....167..176S" ]
[ "astronomy" ]
8
[ "methods: numerical", "methods: observational", "ISM: clouds", "ISM: magnetic fields", "ISM: molecules", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
[ "1953ApJ...118..113C", "1961hhs..book.....C", "1975ApJ...196L..77A", "1983A&A...117..220H", "1983ApJ...270..511Z", "1984ApJ...282..508M", "1987ASIC..210..491M", "1995ApJ...444L.105M", "1996ApJ...460..777F", "1997A&A...325..329T", "1997SoPh..175...93N", "2000ApJS..131..273F", "2004ApJ...603..584P", "2008A&A...486L..13A", "2008ASPC..385..145D", "2008ApJ...680..428G", "2009ApJ...696..567H", "2010A&A...518L.104M", "2010ApJ...716..427B", "2010Natur.466..947B", "2011ApJ...734...63S", "2011ApJ...741...21C", "2011ApJS..192....9T", "2012A&A...543L...3H", "2012ApJ...761L...4B", "2012MNRAS.420.1562H", "2013A&A...550A..38P", "2013A&A...556A..65E", "2013JCoPh.243..269L", "2013MNRAS.436.3707L", "2014A&A...568A..98A", "2014ApJ...789...82C", "2014MNRAS.444.2507P", "2015A&A...575A.110H", "2015ApJ...806..132G", "2015ApJ...807....5F", "2015ApJ...811...71Q", "2015PhRvL.115x1302C", "2016A&A...586A.135P", "2016A&A...591A.104H", "2016ApJ...817...44F", "2016ApJ...820...38H", "2016MNRAS.458..789T", "2016MNRAS.460.1934M", "2016MNRAS.461.3918H", "2016MNRAS.462.1517P" ]
[ "10.1093/mnras/stw1881", "10.48550/arXiv.1607.08615" ]
1607
1607.08615_arXiv.txt
Star formation occurs in condensations located within the dense elongated structures of molecular clouds. These structures, referred to as filaments, have been extensively studied both observationally and theoretically (see review of Andr{\'e} et al. 2014). Although the role of the magnetic field in the evolution of filaments is still a topic of debate, its topology with respect to these filaments is well established. Polarimetric studies (Moneti et al. 1984; Pereyra \& Magalh{\~a}es 2004; Alves et al. 2008; Chapman et al. 2011; Sugitani et al. 2011; Palmeirim et al. 2013; Planck Collaboration et al. 2014) have revealed that the magnetic field is well ordered near dense filaments and perpendicular to their long axis. Elongated structures, called striations, are also seen in the low column density parts of molecular clouds. Despite the fact that striations are not sites of star formation they are of high importance for interstellar medium (ISM) studies since they can reveal the dynamics of molecular clouds, and the early stages of star formation. However, there is yet no theoretically established physical mechanism explaining their formation. Understanding how striations form, whether they are long-lived or transient features and their role in star formation are important open questions. Striations were first observed in $^{12}\rm{CO}$ and $^{13}\rm{CO}$ by Goldsmith et al. (2008) at the northwest part of the Taurus molecular cloud where they appear as autonomous structures. Striations were also observed by $\textit{Herschel}$ in dust emission. One of the most representative examples is the Polaris flare where well ordered, low density elongations are seen throughout the cloud (Miville-Desch{\^e}nes et al. 2010). Like in Taurus, striations in the Polaris flare do not appear to be associated with the denser parts of the cloud. However, in certain clouds, striations are connected to denser filaments. Hennemann et al. (2012), Palmeirim et al. (2013) and Alves de Oliveira et al. (2014) analysed $Herschel$ dust emission maps from DR21, Taurus and Chamaeleon molecular clouds respectively. In all of these studies, striations were interpreted as streamlines in which material flows into or out from more dense filaments and/or clumps. Malinen et al. (2015) compared $Herschel$ dust emission maps and Plank polarization data from the cloud L1642 in order to quantify the relative angle between the plane-of-the-sky component of the magnetic field and the long axes of striations. Using the Rolling Hough Transform (RHT) algorithm (Clark et al. 2014) they concluded that striations were in excellent alignment with the magnetic field. Panopoulou et al. (2016b) also used RHT to compare the orientation of the plane-of-the-sky (POS) magnetic field with linear structures in the Polaris flare and reported that the majority of striations were aligned with the projected magnetic field. The alignment between these structures and the magnetic field has also been pointed out in all clouds by all relevant studies in the literature (Goldsmith et al. 2008; Chapman et al. 2011; Hennemann et al. 2012; Palmeirim et al. 2013; Alves de Oliveira et al. 2014). Li et al. (2013) considered the overall morphology of the magnetic field with respect to both filaments and striations. They concluded that besides the formation of dense filaments from the gravitational contraction along field lines, strong magnetic fields could also act as the guiding channels of sub-\Alf ic flows, thus forming striations. In this mass-accretion/flows-along-field-lines paradigm, which is currently the most common interpretation of such structures, density fluctuations are presumably caused by pressure differences which are in turn caused by fluctuations of the streaming speed as expected by Bernoulli's principle. A shear velocity between ambient streamlines would normally lead to a Kelvin-Helmholtz instability. However, the presence of the magnetic field can stabilize the flow as long as the velocity difference between ambient streamlines is less than two times the \Alf~speed (Frank et al. 1996). In an early theoretical work, Frank et al. (1996) performed 2D simulations in ideal magnetohydrodynamics (MHD) assuming super-\Alf ic velocities with opposite signs on either side of a shear layer and an initially ordered magnetic field. They showed that although a Kelvin-Helmholtz instability occurred early on, a stable, laminar flow was quickly developed due to the presence of the magnetic field. The final density configuration in their simulations was parallel elongated structures aligned with the magnetic field. Supersonic motions and other kinematic properties of molecular clouds have often been interpreted in terms of the presence of hydromagnetic waves (Arons \& Max 1975; Zweibel \& Josafatsson 1983). Specifically, the linewidth-size relation is attributed to \Alf~waves with long wavelengths and large amplitude (Mouschovias \& Psaltis 1995). These findings suggest that striations may also be connected to MHD waves. In the present paper we explore four possible physical mechanisms that could create striations. Since flows along magnetic field lines have been proposed by previous observational studies and are currently considered to be the most plausible mechanism for the formation of striations, we explore two models involving such flows. In the first model, we assume a sub-\Alf ic bulk flow and sub-\Alf ic velocity gradients between ambient streamlines. In the second scenario, we repeat the super-\Alf ic simulations performed by Frank et al. (1996) by adopting values for the parameters involved appropriate for molecular cloud conditions. In the third model, sub-\Alf ic flows perpendicular to the magnetic field cause a Kelvin-Helmholtz instability which in turn produces striations. Finally, we consider an entirely different scenario in which striations are formed from the excitation of compressional magnetosonic waves. In this model, fluctuations of magnetic pressure create striations. Magnetosonic waves are naturally excited from \Alf~waves due to density inhomogeneities. The values adopted in our models are driven from observational results from the literature and analysis of observational data presented here. We find that in the first three models the density contrast between the linear structures that are formed is so low that essentially flows along or perpendicular to magnetic field lines fail to create striations. In contrast, the model including coupling of MHD waves successfully reproduces most of the observational properties of striations. In section \S~\ref{observations} we quantify the observed properties of striations to facilitate a quantitative comparison to simulations. Numerical simulations of models involving streamers and corresponding results are described in \S~\ref{sub1}, \S~\ref{super} and \S~\ref{sub2}. In \S~\ref{ms_from_alf} we provide some theoretical background for our fourth physical model (MHD waves) and describe our results. We summarize and discuss our conclusions in \S~\ref{discuss}.
\label{discuss} The current picture for the formation of striations includes streams that flow along magnetic field lines. We performed numerical simulations adopting two models involving such streamers, a model in which elongated structures are created as a result of a Kelvin-Helmholtz instability perpendicular to field lines and a new model in which striations are formed from the excitation of magnetosonic waves. We assessed the validity of each of our models by comparison between simulated and observational results based on four criteria:\begin{inparaenum}[\itshape a\upshape)] \item the contrast between minima and maxima in density and column density maps \item the spatial power spectrum in each velocity slice and in column density \item the kinematic properties (i.e. velocity range) \item whether the abundance of $\rm{CO}$ follows the total density~\end{inparaenum} and the contrast in abundance is significant for striations to be observed. We proved that flows, either sub-\Alf ic or super-\Alf ic, cannot reproduce the observed contrast even for huge velocity differences between ambient streamlines. The maximum possible contrast in the simulations involving flows is $\sim$ 0.03 \%. The mean contrast in observations is $\sim$ 25 \%. In our second model in which super-\Alf ic streamers flow along magnetic field lines, the contrast is low in both the 2D and 3D models. That is because the thickness of the LOS dimension is constant and thus cancelled out when computing the contrast in column density maps. As a result, projection effects are of minor importance and the observed contrast is an intrinsic density contrast rather than a geometrical effect. Flows perpendicular to field lines is a mechanism that can also qualitatively produce elongated structures. However, due to development of turbulence, these structures are not long-lived. Furthermore, specific projection angles are required so that these structures are seen parallel to the magnetic field. A scenario in which this mechanism can simultaneously produce low density striations parallel to field lines and dense filaments perpendicular to the magnetic field is also difficult to realize. Finally and most importantly, such flows fail to reproduce the observed contrast of striations by more than three orders of magnitude. Overall, this mechanism cannot account for the formation of striations. In the low column density parts of molecular clouds there is good coupling between matter and magnetic field since the degree of ionization is large. Hence, in a paradigm where magnetic field lines act as flux tubes and striations are formed from flows along field lines there must also be regions of stronger and weaker magnetic field. By definition however, such a configuration is equivalent to a wave travelling perpendicular to the long axis of the striations. The quasi-periodicity seen in perpendicular cuts in observations also suggests a formation mechanism which includes superposition of waves. In contrast to streamers, a model including coupling of MHD waves is physical and can naturally explain the formation of striations. Furthermore, for a certain set of parameters the contrast can be up to 7 \%. Besides the large number of combinations (length of LOS dimension, density and magnetic field values) that can be realized and could alter this value, it would certainly be enhanced due to chemical effects. Radiative transfer effects might also be important in an intensity map. Therefore, we conclude that this model can account for the observed contrast. However, even in the 3D simulations performed here the total velocity range over which striations appear is a factor of $\sim$3 smaller compared to observations. There is a number of possibilities to explain this shortcoming. First, intrinsic magnetic field and density values in Taurus could be outside the parameter space considered. From the right panel of Figure~\ref{parameter_space} it can be seen that the lower the plasma $\rm{\beta}$ the larger the velocity range. The amplitude of the perturbation is also a key parameter that affects both the maximum range and the evolution of each velocity component. A second possibility is the existence of multiple sheet-like resonant structures along the LOS which move with respect to each other. In such a picture all of these sheet-like structures should have approximately the same boundaries in order for the same dominant frequency to be present in the spatial power spectrum in all velocity slices. We have demonstrated that when we consider a spectrum of \Alf~waves initially present in the system results are even more realistic and the velocity range increases by more than a factor of two compared to the simple case of having one \Alf~wave. Altering the distribution of power could further increase the velocity range. Although numerical dissipation does not significantly affect the velocity range (see Appendix~\ref{convergence}), the growth of transversal gradients could be affected by the boundary conditions. In a recent paper Hacar et al. (2016) presented a thorough analysis of the $\rm{CO}$ data also used in this paper. They concluded that suprathermal $\rm{CO}$ linewidths could be explained from optical-depth effects and multiple narrow components the superposition of which act as a broadening mechanism. Based on this interpretation of $\rm{CO}$ linewidths they suggested that intrinsic gas motions were transonic. As a result, the velocities found for the majority of the simulations in the parameter study could be within observational ranges. Despite the observational features for which our fourth model can account for, the question arises why we considered incompressible \Alf~waves initially present in the system rather than directly setting up compressible magnetosonic waves. \Alf~waves are exact solutions of the equations of ideal MHD and are thus longlived. Zweibel \& Josafatsson (1983) studied the damping mechanisms of MHD waves and naturally found that \Alf~waves were the longest lived mode. \Alf~waves can also act as the energy carriers from remote regions than the ones where striations are formed and are ultimately observed. Hence, in this model, they arise naturally as the source from which magnetosonic waves pump energy. The spectrum of \Alf~waves passing through an inhomogeneous region is of great importance. The energy distribution in the power spectrum will ultimately be a function of the properties of \Alf~waves initially present in the system. Consequently, the power spectrum of striations could be used to study the spectrum of \Alf~waves present in that region. We intend to return to the problem in follow-up publication with more 3D simulations and a larger parameter space. The effect that different dimensions and projection angles have on the power spectrum is also left for future study. Mouschovias (1987) predicted that torsional \Alf~waves can naturally be generated by the rotation of a clump and can also be trapped between magnetically linked clumps. Just as linear \Alf~waves, in the non-linear regime, torsional \Alf~waves can also excite fast magnetosonic waves (Tirry \& Berghmans 1997). As a result, striations connected to denser filaments could also be explained through the same mechanism. Thus, the interplay between \Alf~and magnetosonic waves along with acoustic waves and gravitational contraction along magnetic field lines is a promising scenario for explaining the overall gas-magnetic field morphology. Additional 3D simulation including gravity, will determine if the phase mixing between torsional \Alf~waves and fast magnetosonic waves can reproduce the observed properties of striations associated with denser parts of molecular clouds. Elongated structures, usually referred to as fibers, have also been observed at high Galactic latitudes in the diffuse interstellar medium (see Clark et al. 2015 and references therein). Similar to striations, the magnetic field in these regions is well ordered and parallel to fibers which again exhibit quasi-periodicity. Thus, it is possible that striations and fibers share a common formation mechanism.
16
7
1607.08615
Dust continuum and molecular observations of the low column density parts of molecular clouds have revealed the presence of elongated structures which appear to be well aligned with the magnetic field. These so-called striations are usually assumed to be streams that flow towards or away from denser regions. We perform ideal magnetohydrodynamic (MHD) simulations adopting four models that could account for the formation of such structures. In the first two models striations are created by velocity gradients between ambient, parallel streamlines along magnetic field lines. In the third model striations are formed as a result of a Kelvin-Helmholtz instability perpendicular to field lines. Finally, in the fourth model striations are formed from the non-linear coupling of MHD waves due to density inhomogeneities. We assess the validity of each scenario by comparing the results from our simulations with previous observational studies and results obtained from the analysis of CO (J = 1-0) observations from the Taurus molecular cloud. We find that the first three models cannot reproduce the density contrast and the properties of the spatial power spectrum of a perpendicular cut to the long axes of striations. We conclude that the non-linear coupling of MHD waves is the most probable formation mechanism of striations.
false
[ "magnetic field lines", "lines", "striations", "elongated structures", "molecular clouds", "molecular observations", "such structures", "denser regions", "MHD waves", "density inhomogeneities", "parallel streamlines", "the magnetic field", "results", "previous observational studies", "velocity gradients", "MHD", "the Taurus molecular cloud", "the fourth model striations", "the low column density parts", "Taurus" ]
11.757286
10.35888
-1
12593739
[ "Beutler, Florian", "Seo, Hee-Jong", "Saito, Shun", "Chuang, Chia-Hsun", "Cuesta, Antonio J.", "Eisenstein, Daniel J.", "Gil-Marín, Héctor", "Grieb, Jan Niklas", "Hand, Nick", "Kitaura, Francisco-Shu", "Modi, Chirag", "Nichol, Robert C.", "Olmstead, Matthew D.", "Percival, Will J.", "Prada, Francisco", "Sánchez, Ariel G.", "Rodriguez-Torres, Sergio", "Ross, Ashley J.", "Ross, Nicholas P.", "Schneider, Donald P.", "Tinker, Jeremy", "Tojeiro, Rita", "Vargas-Magaña, Mariana" ]
2017MNRAS.466.2242B
[ "The clustering of galaxies in the completed SDSS-III Baryon Oscillation Spectroscopic Survey: anisotropic galaxy clustering in Fourier space" ]
298
[ "Institute of Cosmology &amp; Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK; Lawrence Berkeley National Lab, 1 Cyclotron Rd, Berkeley, CA 94720, USA", "Department of Physics and Astronomy, Ohio University, 251B Clippinger Labs, Athens, OH 45701, USA", "Kavli Institute for the Physics and Mathematics of the Universe (WPI), The University of Tokyo Institutes for Advanced Study, The University of Tokyo, Kashiwa, Chiba 277-8583, Japan; Max-Planck-Institut f ür Astrophysik, Karl-Schwarzschild-Strasse 1, D-85740 Garching bei M ünchen, Germany", "Instituto de F ísica Te órica, (UAM/CSIC), Universidad Aut ónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Leibniz-Institut f ür Astrophysik Potsdam (AIP), An der Sternwarte 16, D-14482 Potsdam, Germany", "Institut de Ci ències del Cosmos (ICCUB), Universitat de Barcelona (IEEC-UB), Mart í i Franqu ès 1, E-08028 Barcelona, Spain", "Harvard -Smithsonian Center for Astrophysics, 60 Garden St, Cambridge, MA 02138, USA", "Sorbonne Universit és, Institut Lagrange de Paris (ILP), 98 bis Boulevard Arago, F-75014 Paris, France; Laboratoire de Physique Nuclaire et de Hautes Energies, Universit é Pierre et Marie Curie, 4 Place Jussieu, F-75005 Paris, France", "Max-Planck-Institut f ür extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany; Universit äts-Sternwarte M ünchen, Ludwig-Maximilians-Universit ät M ünchen, Scheinerstra ße 1, D-81679 M ünchen, Germany", "Department of Astronomy, University of California, Berkeley, CA 94720, USA", "Leibniz-Institut f ür Astrophysik Potsdam (AIP), An der Sternwarte 16, D-14482 Potsdam, Germany", "Department of Physics, University of California, Berkeley, CA 94720, USA", "Institute of Cosmology &amp; Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK", "Institute for Gravitation and the Cosmos, The Pennsylvania State University, University Park, PA 16802, USA", "Institute of Cosmology &amp; Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK", "Instituto de F ísica Te órica, (UAM/CSIC), Universidad Aut ónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Campus of International Excellence UAM +CSIC, Cantoblanco, E-28049 Madrid, Spain; Instituto de Astrof ísica de Andaluc ía (CSIC), Glorieta de la Astronom ía, E-18080 Granada, Spain", "Max-Planck-Institut f ür extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85741 Garching, Germany", "Instituto de F ísica Te órica, (UAM/CSIC), Universidad Aut ónoma de Madrid, Cantoblanco, E-28049 Madrid, Spain; Campus of International Excellence UAM +CSIC, Cantoblanco, E-28049 Madrid, Spain", "Institute of Cosmology &amp; Gravitation, Dennis Sciama Building, University of Portsmouth, Portsmouth PO1 3FX, UK; Sorbonne Universit és, Institut Lagrange de Paris (ILP), 98 bis Boulevard Arago, F-75014 Paris, France", "Institute for Astronomy, University of Edinburgh, Royal Observatory, Edinburgh EH9 3HJ, UK", "Department of Astronomy and Astrophysics, The Pennsylvania State University, University Park, PA 16802, USA; Institute for Gravitation and the Cosmos, The Pennsylvania State University, University Park, PA 16802, USA", "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, St Andrews, Fife KY16 9SS, UK", "Instituto de Fisica, Universidad Nacional Autonoma de Mexico, Apdo. Postal 20-364, Mexico City, M éxico." ]
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[ "astronomy" ]
32
[ "gravitation", "surveys", "cosmological parameters", "cosmology: observations", "dark energy", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1979Natur.281..358A", "1987MNRAS.227....1K", "1992StaSc...7..457G", "1994ApJ...426...23F", "1996AJ....111.1748F", "1996ApJ...470L...1M", "1996MNRAS.282..877B", "1998AJ....116.3040G", "2001Natur.410..169P", "2002AJ....123.2121S", "2003ApJ...598..720S", "2003MNRAS.346...78H", "2004PhRvD..70h3007S", "2005ApJ...620..559J", "2006AJ....131.2332G", "2006PASJ...58...93Y", "2006PhR...429..307L", "2006PhRvD..74l3507T", "2007A&A...464..399H", "2007ApJ...664..675E", "2007ApJ...665...14S", "2008ApJ...676..889O", "2008Natur.451..541G", "2008PThPh.120..609Y", "2008PhRvD..77l3540P", "2009ApJ...693.1404S", "2009JCAP...08..020M", "2010AJ....139.1628D", "2010PhRvD..82f3522T", "2011AJ....142...72E", "2011MNRAS.415.2876B", "2011MNRAS.416.3017B", "2012AJ....144..144B", "2012ApJ...756..127G", "2012MNRAS.420.2102S", "2012MNRAS.423.3430B", "2012MNRAS.424..564R", "2012MNRAS.425..405B", "2012MNRAS.426.2719R", "2012PhRvD..86h3540B", "2012PhRvD..86j3528T", "2013AJ....145...10D", "2013AJ....146...32S", "2013MNRAS.429.1514S", "2013MNRAS.433.3559C", "2013PASP..125..306F", "2014A&A...562A..23G", "2014MNRAS.439.2515O", "2014MNRAS.439.2531P", "2014MNRAS.441...24A", "2014MNRAS.441.1457B", "2014MNRAS.443.1065B", "2014MNRAS.444.1400N", "2014MNRAS.444.3501B", "2014PhRvD..90l3522S", "2015ApJ...798....7B", "2015ApJS..219...12A", "2015MNRAS.447.1789Y", "2015MNRAS.453L..11B", "2015PhRvD..92h3532S", "2015PhRvD..92j3516O", "2016A&A...594A..13P", "2016JCAP...02..018L", "2016JCAP...08..050Z", "2016MNRAS.455.1553R", "2016MNRAS.455.3230B", "2016MNRAS.455.4046M", "2016MNRAS.456.4156K", "2016MNRAS.457.4021L", "2016MNRAS.457.4340K", "2016MNRAS.460.1173R", "2016MNRAS.460.1457S", "2016MNRAS.460.3624S", "2016MNRAS.460.4188G", "2016PhRvD..93f3507L", "2016PhRvD..93f3512S", "2017MNRAS.464.3121W", "2018PhRvD..97f3526L" ]
[ "10.1093/mnras/stw3298", "10.48550/arXiv.1607.03150" ]
1607
1607.03150.txt
\label{sec:intro} Clustering in the matter density field carries an enormous amount of information about cosmological parameters. The growth of the matter clustering amplitude is directly sensitive to the sum of the neutrino masses~\citep[e.g.,][]{Lesgourgues:2006nd,Beutler:2014}, the dark energy equation of state and the nature of Gravity~\citep{Kaiser:1987qv,Peacock:2001gs,Guzzo:2008ac}. Galaxy redshift surveys sample the underlying matter density field with galaxies as tracer particles. A measurement of the matter clustering amplitude, $\sigma_8$, could be compared to the precise measurement of the matter clustering amplitude at the recombination redshift, measured in the Cosmic Microwave Background (CMB) providing a long lever-arm with which to test the growth of structure. The formation of galaxies is correlated with the underlying density field, but the galaxy formation processes are complicated and do not allow us to infer these correlations from first principles. I.e., while the clustering amplitude of the galaxy density field can be measured to percent level precision, it cannot straightforwardly be related to the clustering amplitude of the matter density field due to the uncertainties in the bias relation. Despite these limitations, galaxy redshift surveys still allow constraints on the matter clustering amplitude due to redshift-space distortions (RSD). RSD are caused by the underlying peculiar velocity field along the line-of-sight that Doppler-shifts the features in the spectral energy distribution of a galaxy. If the redshift is used to estimate the distance to the galaxy using Hubble's law, peculiar velocity contributions to the redshift introduce errors in the physical coordinates along the line-of-sight. Since it is nearly impossible to estimate the line-of-sight peculiar velocities of millions of galaxies individually, with a precision anywhere near the redshift uncertainty, rather than correcting the effect in the redshift measurements, we account for the resulting statistical distortions in the observed clustering signal. Peculiar velocities trace the gravitational potential field, itself produced by the underlying matter density field. Therefore the distortion effect due to the peculiar velocity field is correlated with the density field. Since the peculiar velocity field affects only the line-of-sight positions without affecting the angular positions, the distortions generate anisotropy in the observed clustering. In the linear regime, redshift-space distortions lead to an angle-dependent increase in the power spectrum amplitude of $(1 + \beta\mu^2)^2$~\citep{Kaiser:1987qv}, where $\mu$ is the cosine of the angle to the line-of-sight and $\beta = f/b$ is the growth rate $f$ divided by the galaxy bias $b$. Using the angular dependence of the RSD signal we can constrain the parameter combination $f(z)\sigma_8(z)$. The redshift-space distortion signal is now considered one of the most powerful observables in large-scale structure~\citep{Peacock:2001gs,Hawkins:2002sg,Tegmark:2006az,Guzzo:2008ac,Yamamoto:2008gr,Blake:2011rj,Beutler:2012px,Reid:2012sw,Samushia:2012iq,Chuang:2013hya,Nishimichi:2013aba}. Another source of anisotropy in the galaxy clustering signal is known as the Alcock-Pazynski (AP) effect~\citep{Alcock:1979mp}. The AP effect is imprinted in the clustering measurement when converting from observable (redshifts and angles) to physical coordinates, which requires a fiducial cosmology. If that fiducial cosmology deviates from the true cosmology, it distorts the clustering signal differently along the line-of-sight and in angular scales. The two effects would be difficult to separate with a featureless power spectrum. By measuring the distortion in the distinct Baryon Acoustic Oscillations (BAO) feature present in the power spectrum, we can constrain the AP signal, thereby isolating the anisotropy in the clustering amplitude due to the RSD~\citep{Matsubara:1996nf,Ballinger:1996cd,Padmanabhan:2008ag,Okumura:2007br}. The AP test through the BAO signal constrains the geometry, i.e., $D_V(z)/r_s(z_d)$ and $F_{\rm AP}(z) = (1+z)D_A(z)H(z)/c$, where $D_V(z) = [(1+z)^2D_A^2(z)cz/H(z)]^{1/3}$ is the angle averaged distance depending on the angular diameter distance $D_A(z)$ and the Hubble parameter $H(z)$. In this paper we use the final data release (\citealt{Alam:2015mbd}, DR12) of the Baryon Oscillation Spectroscopic Survey~\citep{Dawson:2012va}, the largest galaxy redshift dataset available to date to measure the anisotropy in the galaxy power spectrum. Our analysis framework follows our DR11 analysis~\citep{Beutler:2013yhm} with several modifications: (1) In addition to the monopole and quadrupole we now include the hexadecapole, (2) we modified the fitting range to $k = 0.01$ - $0.15\ihMpc$ for the monopole and quadrupole and $k = 0.01$ - $0.10\ihMpc$ for the hexadecapole, (3) we modified our method to include the effect due to the discrete \textit{k}-space grid when estimating the power spectrum, (4) we use a computationally more efficient way to include window function effects, (5) we use larger \textit{k}-bins to reduce the noise in the covariance matrix estimation, (6) we use a new set of mock catalogues called MultiDark-Patchy, which have been introduced in~\citet{Kitaura:2015uqa} and (7) we employ the FFT based power spectrum estimator suggested by~\citet{Bianchi:2015oia} and~\citet{Scoccimarro:2015bla} instead of the $\mathcal{O}(N^2)$ algorithm we used previously to speed up the computation. Our companion paper,~\cite{Beutler:2016}, presents a BAO-only analysis, where we use the same power spectrum multipole measurements, but isolate the BAO signal by marginalising over the shape of the power spectrum (including RSD). While the BAO-only analysis does not capture the information on RSD, it will allow measurements of $D_V(z)/r_s(z_d)$ and $F_{\rm AP}(z) = (1+z)D_A(z)H(z)/c$ that do not depend on our understanding of redshift-space distortions. Without the need to model RSD in detail, the BAO-only analysis can use a wider range of wave numbers ($k_{\rm max} = 0.3 \ihMpc$), and improve the BAO signal by using the BAO reconstruction technique~\citep{Eisenstein:2006nk}. BAO measurements obtained using the monopole and quadrupole correlation function are presented in~\citet{Ross:2016}, while~\citet{Vargas-Magana2016} diagnoses the level of theoretical systematic uncertainty in the BOSS BAO measurements. Beside this paper there are three more measurements of the growth of structure~\citep{Grieb:2016, Sanchez:2016, Satpathy2016}. \citet{Alam2016} combines the results of these seven papers (including this work) into a single likelihood that can be used to test cosmological models. This paper is organised as follows: Section~\ref{sec:data} presents the BOSS dataset used in this analysis. Section~\ref{sec:estimator} describes the power spectrum estimator, followed by the treatment of the window function in section~\ref{sec:win}. Section~\ref{sec:model} introduces our power spectrum model, which is based on renormalised perturbation theory. Section~\ref{sec:mocks} discusses the mock catalogues used to derive covariance matrices for our measurements. In Section~\ref{sec:sys} we use these mock catalogues as well as N-body simulations to test our power spectrum model. The analysis together with the fitting results is presented in section~\ref{sec:analysis}, while section~\ref{sec:dis} discusses the result and compares to previous studies. We conclude in section~\ref{sec:conclusion}. The fiducial cosmological parameters, which we use to turn the observed angles and redshifts into co-moving coordinates and to generate our linear power spectrum models as an input for the power spectrum templates, follow a flat $\Lambda$CDM model with $\Omega_m=0.31$, $\Omega_bh^2=0.022$, $h=0.676$, $\sigma_8=0.8$, $n_s=0.96$, $\sum m_{\nu}=0.06\,$eV and $r_s^{\rm fid}(z_d) = 147.78\,$Mpc.
\label{sec:conclusion} We measure the power spectrum multipoles from the final BOSS DR12 dataset in the redshift range $0.2 < z < 0.75$. We extract the Baryon Acoustic Oscillation, Alcock-Paczynski and redshift-space distortion signals using a model based on renormalised perturbation theory. For the first time we include the hexadecapole in our analysis, while appropriately accounting for the survey window function, which reduced the uncertainties on $f(z)\sigma_8(z)$ by about $20\%$. The main results of this analysis are: \begin{enumerate} \item A FFT based window function method, first suggested in~\citet{Wilson:2015lup} using the global plane parallel approximation, can be derived within the local plane parallel approximation, which makes it applicable to wide-angle surveys like BOSS. We present the detailed derivation in appendix~\ref{app:window}. \item By fitting the monopole and quadrupole between $k = 0.01$ - $0.15\ihMpc$ and the hexadecapole between $k = 0.01$ - $0.10\ihMpc$, we were able to extract the constraints $f(\zeff)\sigma_8(\zeff) = 0.482\pm0.053$, $0.455\pm0.050$ and $0.410\pm0.042$ at the effective redshifts of $z_{\rm eff} = 0.38, 0.51$, and $0.61$, respectively. For the Alcock-Paczynski parameter $F_{\rm AP} = (1+\zeff)D_A(\zeff)H(\zeff)/c$ we find $0.427\pm0.022, 0.594\pm0.035$, and $0.736\pm0.040$ and the BAO scale parameter is $D_Vr_s^{\rm fid}/r_s = 1493\pm28, 1913\pm35$, and $2133\pm36\,$Mpc. Assuming Gaussian likelihood, we provide a covariance matrix which contains the parameter constraints as well as their correlations (see appendix~\ref{app:covs}). \item We demonstrated the accuracy of our analysis pipeline by participating in a mock challenge, which resulted in systematic uncertainties $\lesssim 10\%$ of the statistical error budget. The description of this mock challenge can be found in our companion paper~\citep{Tinker:2015}. \item Our high redshift result on $f\sigma_8$ is in agreement with our DR11 analysis using the CMASS sample, and shows a small $1.4\sigma$ deviation from the Planck prediction. The low redshift results obtained in this analysis show good agreement with the Planck prediction. \end{enumerate} \citet{Alam2016} combines our measurements with the corresponding growth of structure measurements of~\citet{Grieb:2016},~\citet{Sanchez:2016} and~\citet{Satpathy2016} and the BAO-only measurements of~\citet{Beutler:2016} and~\citet{Ross:2016} into a final BOSS likelihood and investigates the cosmological implications.
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7
1607.03150
We investigate the anisotropic clustering of the Baryon Oscillation Spectroscopic Survey (BOSS) Data Release 12 sample, which consists of 1198 006 galaxies in the redshift range 0.2 &lt; z &lt; 0.75 and a sky coverage of 10 252 deg<SUP>2</SUP>. We analyse this data set in Fourier space, using the power-spectrum multipoles to measure redshift-space distortions simultaneously with the Alcock-Paczynski effect and the baryon acoustic oscillation scale. We include the power-spectrum monopole, quadrupole and hexadecapole in our analysis and compare our measurements with a perturbation-theory-based model, while properly accounting for the survey window function. To evaluate the reliability of our analysis pipeline, we participate in a mock challenge, which results in systematic uncertainties significantly smaller than the statistical uncertainties. While the high-redshift constraint on fσ<SUB>8</SUB> at z<SUB>eff</SUB> = 0.61 indicates a small (∼1.4σ) deviation from the prediction of the Planck ΛCDM (Λ cold dark matter) model, the low-redshift constraint is in good agreement with Planck ΛCDM. 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 in Alam et al. to produce the final cosmological constraints from BOSS.
false
[ "systematic uncertainties", "Planck ΛCDM", "BOSS", "model", "good agreement", "Fourier space", "lt", "z", "the statistical uncertainties", "ΛCDM", "redshift-space distortions", "the baryon acoustic oscillation scale", "the final galaxy clustering data", "the final cosmological constraints", "the survey window function", "the redshift range", "the low-redshift constraint", "z &lt", "the high-redshift constraint", "hexadecapole" ]
12.158066
2.616943
161
17811078
[ "Macciò, Andrea V.", "Udrescu, Silviu M.", "Dutton, Aaron A.", "Obreja, Aura", "Wang, Liang", "Stinson, Greg R.", "Kang, Xi" ]
2016arXiv160701028M
[ "NIHAO X: Reconciling the local galaxy velocity function with Cold Dark Matter via mock HI observations" ]
3
[ "-", "-", "-", "-", "-", "-", "-" ]
[ "2016ApJ...827L..36B", "2016ApJ...832...11B", "2017arXiv170401832S" ]
[ "astronomy" ]
7
[ "Astrophysics - Astrophysics of Galaxies" ]
[ "1989MNRAS.237.1127C", "1999ApJ...522...82K", "1999ApJ...524L..19M", "1999PhDT........27S", "2000ApJ...539..517B", "2001ApJS..137....1B", "2001PhDT........21S", "2001cghr.confE..64H", "2002MNRAS.333..156B", "2003PASP..115..763C", "2004NewA....9..137W", "2005Natur.435..629S", "2006MNRAS.373.1074S", "2009ApJ...700.1779Z", "2009IEEEP..97.1482D", "2010MNRAS.402.1995M", "2010MNRAS.403.1969Z", "2010MNRAS.407.1581S", "2011ApJ...739...38P", "2011ApJ...740..102K", "2011ApJ...742...16T", "2011MNRAS.412.1943P", "2011MNRAS.417.1260F", "2013AJ....145..101K", "2013MNRAS.428..129S", "2013MNRAS.430.2427R", "2014A&A...571A..16P", "2014MNRAS.441.2986D", "2014MNRAS.441L...6S", "2014MNRAS.442..176P", "2014MNRAS.442.3013K", "2015MNRAS.453.2133B", "2015MNRAS.454...83W", "2015MNRAS.454.1105S", "2015MNRAS.454.1798K", "2016ApJ...832...11B", "2016MNRAS.455.3841B", "2016MNRAS.456.3542T", "2016MNRAS.457.1931S", "2016MNRAS.457L..74D", "2016MNRAS.458..912V", "2016MNRAS.459..467O", "2017MNRAS.464.2796G", "2017MNRAS.466.4858W" ]
[ "10.48550/arXiv.1607.01028" ]
1607
1607.01028_arXiv.txt
The Cold Dark Matter (CDM) scenario is very successful in reproducing the large scale structure of the Universe \citep{Springel2005} and the anisotropies in the Cosmic Microwave Background \citep{Planck2014}. In a CDM universe structure formation proceeds in a bottom-up fashion, with small structures collapsing earlier and then merging to form larger and larger haloes. This scenario predicts a larger number of low mass structures compared to more massive ones, or in other words a steeply declining halo mass function. Such a prediction seems to be at odds with observational data both around galaxies, e.g. the satellite abundance \citep{Klypin1999, Moore1999} and in the local volume \citep{Cole1989, Zavala2009, Trujillo-Gomez2011}. In the last years several solutions have been proposed to reconcile the observed number of satellites around the Milky Way and Andromeda galaxies with predictions from CDM \citep[][and references therein]{Bullock2000, Benson2002, Maccio2010, Font2011}. All these solutions are based (to a different extent) on the fact that galaxy formation becomes quite inefficient in low mass dark matter haloes, due to ionizing background, Super Novae (SN) explosions and gas removal due to ram pressure. As a result, by taking into account baryonic processes it is possible to bring into agreement the abundance of satellites in the Local Group and CDM predictions \citep[e.g.,][]{Sawala2016}. A more persistent challenge is provided by field galaxies in the local volume and their velocity function. The galaxy velocity function (GVF) is defined as the abundance of galaxies with a given circular velocity and is conceptually similar to the galaxy luminosity function. The advantage is that, for a given cosmological model, circular velocities can nominally be predicted much more accurately than luminosities, since they are less dependent on the still poorly understood physics of galaxy formation. This makes velocity functions a very useful tool to test theoretical models \citep{Zavala2009, Papastergis2011}. Recently \citet[][hereafter K15]{Klypin2015} have presented new measurements of the GVF in the Local Volume ($ D \approx 10$ Mpc) and they have shown that while CDM provides very good estimates of the number of galaxies with circular velocities around and above $70\,\kms$, it fails quite dramatically at lower circular velocities, overestimating by a factor up to five the number of dwarf galaxies in the velocity range $30-50\,\kms$. As pointed out by K15 and other authors \citep[e.g.,][]{Zavala2009} galaxies in this velocity range are practically insensitive to the ionization background and, by not being satellites, they are not affected by gas depletion via ram pressure or by stellar stripping. This makes the mismatch between the observed GVF and the CDM predictions a quite serious problem for the current cosmological model. Moreover, simple modifications to the CDM paradigm, for example by introducing a warm dark matter component \citep[e.g.,][]{Schneider2014}, seem also to be not able to fix this issue, as shown by K15. \citet{Brook2015} showed that a population of galaxies with mass profiles modified by baryonic feedback \citep{DiCintio2014b} is able to explain the GVF significantly better than a model in which a universal cuspy density profile is assumed. In this Letter we revise the effect of galaxy formation on the GVF using the hydrodynamical simulations from the NIHAO project \citep{Wang2015}. In contrast to many previous theoretical studies, which use proxies for \hi linewidths, we directly measure them. The combination of high spatial resolution and large sample size make NIHAO ideally suited to study the relation between the (rotation) velocity inferred by the kinematics of the \hi gas component and the maximum circular velocity of the halo, and thus to shed light on the reasons behind the large discrepancy between the CDM-based predicted and the observed GVF. \vspace{-0.5cm}
\label{sec:conc} The abundance of galaxies with a given velocity linewidth ($V_{50}\equiv W_{50}/2$) provides a very powerful tool to test against observations different models for Dark Matter, from Cold (CDM) to Warm (WDM). The slope of this velocity function at low velocities ($V_{50}<100\,\kms$) departs significantly from the expectations of the standard CDM model, leading to a difference in abundance of about a factor $\approx 8$ at $V_{50} \sim 50$ \kms \citep[][ALFALFA survey]{Papastergis2011}. Such a discrepancy is still in place when deeper observations of the local ($D<10$ Mpc) Universe are used \citep[K15, based on the survey from][]{Karachentsev2013}; moreover a simple cut in the matter power spectrum, as in WDM models, does not provide a viable solution (K15). These kinds of analyses suffer from the fact that \hi discs in low mass galaxies might not be extended enough to probe the full amplitude of the rotation curve \citep{Brook2016}. We reinvestigated this problem using the NIHAO simulation suite to study in detail the relation between the maximum circular velocity of Dark Matter haloes, \Vmax, and the \hi velocity. We found that at high masses the \hi rotation curves are extended enough to fully capture the dynamical velocity of the dark matter, and hence the linewidth based velocities, $V_{50}$, can be directly compared with velocities coming from dissipationless (N-body) simulations. At lower masses, namely $\Vmax < 100\,\kms$, there is a clear bias between the circular velocity of the halo and the width of the \hi velocity distribution. Such a bias is due to both the \hi distribution not being extended enough to reach the peak of the circular velocity and to the departure of the \hi dynamics from a purely rotating disc, due to substantial dispersion in the vertical direction. When such a bias is taken into account, the simulated \hi velocity function is able to match the observed ones well, showing that the CDM model is not intrinsically flawed, but that attempts to realistically ``observe'' fully hydrodynamical simulations might alleviate apparent tensions between observations and theory predictions. In conclusions any comparisons between pure N-body simulations and observations must be taken with a grain of salt. \vspace{-0.5cm}
16
7
1607.01028
We used 87 high resolution hydrodynamical cosmological simulations from the NIHAO suite to investigate the relation between the maximum circular velocity (Vmax) of a dark matter halo in a collisionless simulation and the velocity width of the HI gas in the same halo in the hydrodynamical simulation. These two quantities are normally used to compare theoretical and observational velocity functions and have led to a possible discrepancy between observations and predictions based on the Cold Dark Matter (CDM) model. We show that below 100 km/s, there is clear bias between HI based velocities and Vmax, that leads to an underestimation of the actual circular velocity of the halo. When this bias is taken into account the CDM model has no trouble in reproducing the observed velocity function and no lack of low velocity galaxies is actually present. Our simulations also reproduce the linewidth - stellar mass (Tully-Fisher) relation and HI sizes, indicating that the HI gas in our simulations is as extended as observed. The physical reason for the lower than expected linewidths is that, in contrast to high mass galaxies, low mass galaxies no longer have extended thin HI rotating disks, as is commonly assumed.
false
[ "HI based velocities", "low velocity galaxies", "thin HI", "low mass galaxies", "high mass galaxies", "the maximum circular velocity", "the actual circular velocity", "the observed velocity function", "theoretical and observational velocity functions", "87 high resolution hydrodynamical cosmological simulations", "Vmax", "the velocity width", "CDM", "the HI gas", "the hydrodynamical simulation", "a collisionless simulation", "a dark matter halo", "disks", "NIHAO", "predictions" ]
10.907444
5.752632
-1
12409581
[ "Nesvorný, David", "Vokrouhlický, David", "Roig, Fernando" ]
2016ApJ...827L..35N
[ "The Orbital Distribution of Trans-Neptunian Objects Beyond 50 au" ]
42
[ "Department of Space Studies, Southwest Research Institute, 1050 Walnut St., Suite 300, Boulder, CO 80302, USA", "Institute of Astronomy, Charles University, V Holešovičkách 2, CZ-18000 Prague 8, Czech Republic", "Observatório Nacional, Rua Gal. Jose Cristino 77, Rio de Janeiro, RJ 20921-400, Brazil" ]
[ "2016AJ....152..133K", "2017AJ....154..171P", "2017Ap&SS.362...11I", "2017ApJ...845...27N", "2018AJ....155...75S", "2018AJ....155..260V", "2018AJ....156...33Y", "2018AJ....156..108M", "2018ARA&A..56..137N", "2018ARA&A..56..357S", "2018ApJ...855L...6H", "2018FrASS...5...14L", "2018MNRAS.480.1870S", "2018arXiv180807446R", "2019AJ....157..181V", "2019AJ....157..253L", "2019AJ....158...43K", "2019AJ....158..214C", "2019Icar..331...49N", "2020tnss.book...61K", "2021AJ....162...27C", "2021AcAau.178..257A", "2021ApJ...920L...9A", "2021Icar..35814184W", "2021MNRAS.506..633D", "2021PSJ.....2...14S", "2022A&A...662L...4E", "2022ApJ...938L..23H", "2022ApJS..258...41B", "2022arXiv220600010K", "2022arXiv221016354F", "2023AJ....166..118L", "2023ApJ...947L...4P", "2023FrST....354360M", "2023Icar..40615738N", "2023Icar..40615763W", "2023MNRAS.525..805M", "2023PSJ.....4..145B", "2024ApJ...962L..33H", "2024ApJ...968...47V", "2024Icar..41516057K", "2024PSJ.....5..135G" ]
[ "astronomy" ]
4
[ "Kuiper belt: general", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1994Icar..108...18L", "1997Natur.387..573L", "2003Icar..161..404G", "2004Icar..170..492G", "2008Icar..196..258L", "2008ssbn.book...43G", "2011AJ....142..131P", "2011ApJ...738...13B", "2012AJ....144..117N", "2012ApJ...744L...3B", "2013ApJ...768...45N", "2014A&A...564A..44B", "2014Natur.507..471T", "2015AJ....149..202P", "2015AJ....150...68N", "2015AJ....150...73N", "2015ApJ...806..143V", "2016AJ....151...22B", "2016AJ....151...31S", "2016AJ....152...23V", "2016AJ....152..133K", "2016ApJ...825...94N", "2016ApJ...825L..13S", "2017AJ....153...33L" ]
[ "10.3847/2041-8205/827/2/L35", "10.48550/arXiv.1607.08279" ]
1607
1607.08279_arXiv.txt
In our previous work, we developed a numerical model of Neptune's migration into an outer planetesimal disk (Nesvorn\'y 2015a,b; Nesvorn\'y \& Vokrouhlick\'y 2016; hereafter NV16). By comparing the model results with the observed distribution of Kuiper belt orbits with $a<50$ au (e.g., Petit et al. 2011), we inferred that Neptune's migration must have been long-range, slow and grainy. Here we use the same model to discuss the orbital structure of the Kuiper belt beyond 50 au. We find that objects scattered by Neptune to $a>50$~au are often trapped into mean motion resonances with Neptune which act to raise the perihelion distances to $q>40$ au, and detach the orbits from Neptune. The objects are subsequently released from resonances as Neptune migrates toward its present orbit. The orbital structure of the detached disk with $a>50$ au and $q>40$ au is thus expected to be clustered near Neptune's resonances. Similar results were recently reported in an independent work (Kaib \& Sheppard 2016). Section 2 briefly reviews the numerical method. The results are presented and compared with observations in Section 3. Our conclusions are given in Section~4.
Our simulations with slow migration of Neptune (as required from the inclination constraint; Nesvorn\'y 2015a) lead to the formation of a prominent detached disk with substantial populations of objects concentrated at various MMRs with Neptune. This is an important prediction of the model, which is testable by observations. The current surveys are only starting to have a sufficient sensitivity to probe the orbital distribution of bodies with large perihelion distances (e.g., Shankman et al. 2016). Sheppard et al. (2016) reported several new objects in the detached disk between 50 and 100 au. They found that these objects are near Neptune's MMRs and have significant inclinations ($i>20^\circ$). Interestingly, these findings are consistent with the predictions of our model with slow migration of Neptune. The population census of near-resonant SDOs inferred from observations is also consistent with the model. Our results imply that the detached population at 50-100 au should be $\simeq$5 times larger than the scattering population in the same semimajor axis range, which may have important implications for the origin of Jupiter-family comets. In addition, there seems to be a large population of objects with $q\simeq35$-40 au in the 5:1 MMR (Pike et al. 2015), which cannot be easily explained by the resonant sticking of scattering objects (Yu et al. 2015). Instead, we find it possible that these objects are the low-$q$, easier-to-detect part of the resonant populations that continue to $q>40$ au.
16
7
1607.08279
The dynamical structure of the Kuiper Belt beyond 50 au is not well understood. Here we report results of a numerical model with long-range, slow, and grainy migration of Neptune. The model implies that bodies scattered outward by Neptune to semimajor axes a\gt 50 {au} often evolve into resonances which subsequently act to raise the perihelion distances of orbits to q\gt 40 {au}. The implication of the model is that the orbits with 50\lt a\lt 100 {au} and q\gt 40 {au} should cluster near (but not in) the resonances with Neptune (3:1 at a = 62.6 au, 4:1 at a=75.9 {au}, 5:1 at a=88.0 {au}, etc.). The recent detection of several distant Kuiper Belt Objects (KBOs) near resonances is consistent with this prediction, but it is not yet clear whether the orbits are really non-resonant as our model predicts. We estimate from the model that there should presently be ∼1600-2400 bodies at the 3:1 resonance and ∼1000-1400 bodies at the 4:1 resonance (for q\gt 40 {au} and diameters D\gt 100 km). These results favorably compare with the population census of distant KBOs inferred from existing observations.
false
[ "au", "a=88.0", "resonances", "Neptune", "orbits", "existing observations", "several distant Kuiper Belt Objects", "semimajor axes", "Kuiper Belt Objects", "grainy migration", "distant KBOs", "bodies", "a=75.9", "50 au", "a\\gt", "D\\gt", "q\\gt", "a = 62.6 au", "the resonances", "KBOs" ]
7.77521
16.055288
93
1578621
[ "Ostrovski, Fernanda", "McMahon, Richard G.", "Connolly, Andrew J.", "Lemon, Cameron A.", "Auger, Matthew W.", "Banerji, Manda", "Hung, Johnathan M.", "Koposov, Sergey E.", "Lidman, Christopher E.", "Reed, Sophie L.", "Allam, Sahar", "Benoit-Lévy, Aurélien", "Bertin, Emmanuel", "Brooks, David", "Buckley-Geer, Elizabeth", "Carnero Rosell, Aurelio", "Carrasco Kind, Matias", "Carretero, Jorge", "Cunha, Carlos E.", "da Costa, Luiz N.", "Desai, Shantanu", "Diehl, H. Thomas", "Dietrich, Jörg P.", "Evrard, August E.", "Finley, David A.", "Flaugher, Brenna", "Fosalba, Pablo", "Frieman, Josh", "Gerdes, David W.", "Goldstein, Daniel A.", "Gruen, Daniel", "Gruendl, Robert A.", "Gutierrez, Gaston", "Honscheid, Klaus", "James, David J.", "Kuehn, Kyler", "Kuropatkin, Nikolay", "Lima, Marcos", "Lin, Huan", "Maia, Marcio A. G.", "Marshall, Jennifer L.", "Martini, Paul", "Melchior, Peter", "Miquel, Ramon", "Ogando, Ricardo", "Plazas Malagón, Andrés", "Reil, Kevin", "Romer, Kathy", "Sanchez, Eusebio", "Santiago, Basilio", "Scarpine, Vic", "Sevilla-Noarbe, Ignacio", "Soares-Santos, Marcelle", "Sobreira, Flavia", "Suchyta, Eric", "Tarle, Gregory", "Thomas, Daniel", "Tucker, Douglas L.", "Walker, Alistair R." ]
2017MNRAS.465.4325O
[ "VDES J2325-5229 a z = 2.7 gravitationally lensed quasar discovered using morphology-independent supervised machine learning" ]
68
[ "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; CAPES Foundation, Ministry of Education of Brazil, Brasília - DF 70040-020, Brazil", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; Department of Astronomy, University of Washington, Seattle, WA 98195, USA", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; DAMTP, Centre for Mathematical Sciences, Wilberforce Road, Cambridge CB3 0WA, UK", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "School of Physics, University of Wollongong, Wollongong, NSW 2522, Australia; Australian Astronomical Observatory, North Ryde, NSW, Australia", "Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK; Kavli Institute for Cosmology, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "CNRS, UMR 7095, Institut d'Astrophysique de Paris, F-75014, Paris, France; Department of Physics &amp; Astronomy, University College London, Gower Street, London WC1E 6BT, UK; Sorbonne Universités, UPMC Univ Paris 06, UMR 7095, Institut d'Astrophysique de Paris, F-75014, Paris, France", "CNRS, UMR 7095, Institut d'Astrophysique de Paris, F-75014, Paris, France; Sorbonne Universités, UPMC Univ Paris 06, UMR 7095, Institut d'Astrophysique de Paris, F-75014, Paris, France", "Department of Physics &amp; Astronomy, University College London, Gower Street, London WC1E 6BT, UK", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil; Observatório Nacional, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil", "Department of Astronomy, University of Illinois, 1002 W. Green Street, Urbana, IL 61801, USA; National Center for Supercomputing Applications, 1205 West Clark St., Urbana, IL 61801, USA", "Institut de Ciències de l'Espai, IEEC-CSIC, Campus UAB, Carrer de Can Magrans, s/n, E-08193 Bellaterra, Barcelona, Spain; Institut de Física d'Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, E-08193 Bellaterra (Barcelona), Spain", "Kavli Institute for Particle Astrophysics &amp; Cosmology, PO Box 2450, Stanford University, Stanford, CA 94305, USA", "Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil; Observatório Nacional, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil", "Excellence Cluster Universe, Boltzmannstr. 2, D-85748 Garching, Germany; Faculty of Physics, Ludwig-Maximilians-Universität, Scheinerstr. 1, D-81679 Munich, Germany", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Excellence Cluster Universe, Boltzmannstr. 2, D-85748 Garching, Germany; Faculty of Physics, Ludwig-Maximilians-Universität, Scheinerstr. 1, D-81679 Munich, Germany", "Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA; Department of Physics, University of Michigan, Ann Arbor, MI 48109, USA", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Institut de Ciències de l'Espai, IEEC-CSIC, Campus UAB, Carrer de Can Magrans, s/n, E-08193 Bellaterra, Barcelona, Spain", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA; Kavli Institute for Cosmological Physics, University of Chicago, Chicago, IL 60637, USA", "Department of Physics, University of Michigan, Ann Arbor, MI 48109, USA", "Department of Astronomy, University of California, Berkeley, 501 Campbell Hall, Berkeley, CA 94720, USA; Lawrence Berkeley National Laboratory, 1 Cyclotron Road, Berkeley, CA 94720, USA", "Kavli Institute for Particle Astrophysics &amp; Cosmology, PO Box 2450, Stanford University, Stanford, CA 94305, USA; SLAC National Accelerator Laboratory, Menlo Park, CA 94025, USA", "Department of Astronomy, University of Illinois, 1002 W. Green Street, Urbana, IL 61801, USA; National Center for Supercomputing Applications, 1205 West Clark St., Urbana, IL 61801, USA", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Center for Cosmology and Astro-Particle Physics, The Ohio State University, Columbus, OH 43210, USA; Department of Physics, The Ohio State University, Columbus, OH 43210, USA", "Cerro Tololo Inter-American Observatory, National Optical Astronomy Observatory, Casilla 603, La Serena, Chile", "Australian Astronomical Observatory, North Ryde, NSW 2113, Australia", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil; Departamento de Física Matemática, Instituto de Física, Universidade de São Paulo, CP 66318, CEP 05314-970, São Paulo, SP, Brazil", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil; Observatório Nacional, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil", "George P. and Cynthia Woods Mitchell Institute for Fundamental Physics and Astronomy, and Department of Physics and Astronomy, Texas A&amp;M University, College Station, TX 77843, USA", "Center for Cosmology and Astro-Particle Physics, The Ohio State University, Columbus, OH 43210, USA; Department of Astronomy, The Ohio State University, Columbus, OH 43210, USA", "Department of Astrophysical Sciences, Princeton University, Peyton Hall, Princeton, NJ 08544, USA", "Institut de Física d'Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, E-08193 Bellaterra (Barcelona), Spain; Institució Catalana de Recerca i Estudis Avançats, E-08010 Barcelona, Spain", "Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil; Observatório Nacional, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Dr., Pasadena, CA 91109, USA", "SLAC National Accelerator Laboratory, Menlo Park, CA 94025, USA", "Department of Physics and Astronomy, Pevensey Building, University of Sussex, Brighton BN1 9QH, UK", "Centro de Investigaciones Energéticas, Medioambientales y Tecnológicas (CIEMAT), Madrid, Spain", "Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil; Instituto de Física, UFRGS, Caixa Postal 15051, Porto Alegre, RS - 91501-970, Brazil", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Centro de Investigaciones Energéticas, Medioambientales y Tecnológicas (CIEMAT), Madrid, Spain", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil; ICTP South American Institute for Fundamental Research Instituto de Física Teórica, Universidade Estadual Paulista, São Paulo, Brazil", "Computer Science and Mathematics Division, Oak Ridge National Laboratory, Oak Ridge, TN 37831, USA", "Department of Physics, University of Michigan, Ann Arbor, MI 48109, USA", "Institute of Cosmology &amp; Gravitation, University of Portsmouth, Portsmouth PO1 3FX, UK", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL 60510, USA", "Cerro Tololo Inter-American Observatory, National Optical Astronomy Observatory, Casilla 603, La Serena, Chile" ]
[ "2017AJ....153..219S", "2017ApJ...838L..15L", "2017ApJS..232...15D", "2017MNRAS.466.3088W", "2017MNRAS.471..167J", "2017MNRAS.471.2013A", "2017MNRAS.471.2224N", "2017MNRAS.471.3378H", "2017MNRAS.472...90D", "2017MNRAS.472.4038A", "2017MNRAS.472.5023L", "2017SSRv..212.1743D", "2017arXiv170609424C", "2018A&A...609A..71C", "2018A&A...618A..56D", "2018IAUS..336...80S", "2018MNRAS.473L.116O", "2018MNRAS.474.3391A", "2018MNRAS.476..663K", "2018MNRAS.477L..70W", "2018MNRAS.479.5060L", "2018MNRAS.480.5017A", "2018MNRAS.481.1041T", "2018MNRAS.481.1115C", "2018PASJ...70S..39L", "2018PDU....22..189B", "2019A&A...625A..56M", "2019A&A...625A.119M", "2019MNRAS.483.5649S", "2019MNRAS.484.4726B", "2019MNRAS.486.4987R", "2019MNRAS.489.2525A", "2019MNRAS.489.4741S", "2019MNRAS.490.1743C", "2019arXiv191103867M", "2019arXiv191208977K", "2020A&A...636A..87C", "2020A&A...640A..88C", "2020A&A...642A.193M", "2020MNRAS.491.1408M", "2020MNRAS.494.1308N", "2020MNRAS.494.3491L", "2020MNRAS.494.3750C", "2020MNRAS.494.6072S", "2020MNRAS.497..556H", "2020MNRAS.498.4021U", "2020WDMKD..10.1349F", "2021A&A...646A.126S", "2021A&A...655A.114C", "2021ApJ...921...42S", "2021ApJ...923...16L", "2021MNRAS.501..269D", "2021MNRAS.502.4617T", "2021MNRAS.503.2229M", "2021MNRAS.507.1937B", "2022A&ARv..30....8T", "2022ApJ...927...45C", "2022ApJS..259...27O", "2022KosNT..28e..27K", "2022MNRAS.509..738D", "2022MNRAS.509.2269S", "2022MNRAS.510.3849U", "2023A&A...671A.147S", "2023AJ....165...26J", "2023MNRAS.518.1260S", "2023MNRAS.519.2528M", "2024MNRAS.528.4188B", "2024SSRv..220...23L" ]
[ "astronomy" ]
9
[ "gravitational lensing: strong", "methods: observational", "methods: statistical", "quasars: general", "Astrophysics - Astrophysics of Galaxies" ]
[ "1979Natur.279..381W", "1980ApJS...43..305K", "1980ApJS...43..393C", "1984Msngr..38....9B", "1991ApJS...77..119M", "1992ApJS...79....1T", "1993ASPC...52..173T", "1994ApJ...429..582C", "1994MNRAS.268..305H", "1996A&AS..117..393B", "1998ApJ...509..561K", "1998MNRAS.295..587M", "1999MNRAS.307..225K", "1999MNRAS.310..540A", "2001ApJ...547...50K", "2001MNRAS.322L..29C", "2001astro.ph..2340K", "2002AJ....123..485S", "2002AJ....123.2945R", "2003AJ....126.2281J", "2003ApJ...587..143R", "2003MNRAS.341....1M", "2003MNRAS.341...13B", "2004AJ....127.1318P", "2004ApJ...610...69K", "2006A&A...451..865C", "2006A&A...457..841I", "2006AJ....132..999O", "2006ApJ...637L..73K", "2006ApJ...649..616P", "2008AJ....135..496I", "2008AJ....135..512O", "2008ApJ...673...34P", "2008MNRAS.387..741J", "2008SerAJ.176....1I", "2009ApJS..182..543A", "2010AJ....140..370M", "2010AJ....140..403I", "2010AJ....140.1868W", "2010ApJ...711..201S", "2010ApJ...712.1129M", "2010ApJ...724..511A", "2010MNRAS.405.2579O", "2010PASJ...62.1017O", "2010ascl.soft10012O", "2011ApJ...735..112J", "2011MNRAS.411L...6A", "2012AJ....143..119I", "2012ApJ...753...30S", "2012ApJ...753..106M", "2012Natur.481..341V", "2013A&A...558A..33A", "2013AJ....145...48M", "2013ApJ...766...70S", "2013ApJ...772...26A", "2013Msngr.154...35M", "2014ApJ...792L..19R", "2014ApJ...794L..20M", "2014MNRAS.442.2434N", "2015A&A...579A..40S", "2015AJ....150..150F", "2015ApJ...807...50B", "2015ApJS..219...39R", "2015MNRAS.452.3047Y", "2015MNRAS.452.3124D", "2015MNRAS.454.1260A", "2016A&A...592A...1P", "2016MNRAS.456.1595M", "2016MNRAS.460.1270D", "2016PASA...33....1L", "2017MNRAS.465.4914B" ]
[ "10.1093/mnras/stw2958", "10.48550/arXiv.1607.01391" ]
1607
1607.01391_arXiv.txt
The discovery of the first strong gravitational lens \citep{Walsh+1979} brought forth a powerful tool to study cosmology and astrophysics. Systems where the background source is a quasar can be used to map the dark matter substructure \citep[e.g.][]{MaoSchneider1998,Kochanek+2004,Vegetti+2012,Nierenberg+2014}; to determine the mass \citep[e.g.][]{Morgan+2010} and spin \citep{Reynolds+2014} of black holes; to measure the properties of distant host galaxies \citep[e.g.][]{Kochanek+2001,Claeskens+2006,Peng+2006} and to measure the value of the Hubble constant $H_{0}$. The constraints on cosmological parameters in particular, are comparable in precision to baryonic acoustic oscillation methods \citep[e.g][]{Suyu+2010,Suyu+2013,Bonvin+2016}. In addition to that, the effects of microlensing of the quasar induced by the stars in the lens galaxy can be used to probe the physical properties of quasar accretion disks such as the wavelength dependence of the size of the accretion disk \citep[e.g.][]{Poindexter+2008}. Large samples of new quasar lens systems were discovered through dedicated surveys in radio, like the Cosmic Lens All Sky Survey (CLASS; \citealt{Myers+2003,Browne+2003}) that, in combination with the Jodrell Bank VLA Astrometic Survey (JVAS; \citealt{King+1999}), found 21 new lens systems, and in the optical, such as the SDSS Quasar Lens Search (SQLS; \citealt{Oguri+2006,Oguri+2008,Inada+2008,Inada+2010,Inada+2012}), that discovered 49 new lensed systems\footnote{\url{http://www-utap.phys.s.u-tokyo.ac.jp/~sdss/sqls/lens.html}}. The future of the field is currently hindered by the need for more lenses. This can be achieved, according to \citealt{OguriMarshall2010} (hereafter OM10), not by increasing the depth of the searches, but the area. Therefore, surveys such as the Dark Energy Survey (DES - \citealt{DES2005,Abbott+2016}) and, in the future, the Large Synoptic Survey Telescope (LSST - \citealt{Ivezic+2008}) are capable of more than doubling the current quasar lens sample size. \begin{figure} \begin{center} \includegraphics[width=1.0\columnwidth]{figs/om10predict2.png} \end{center} \caption{Expected number of quasar lenses in the full DES area as function of the magnitude of the brightest quasar image. The blue solid line shows the overall expected number of systems, while the dashed pink line and the dotted green line show the number of pairs and quads, respectively. The dot-dashed black line shows the expected number of lenses as a function of the magnitude of the second brightest quasar image for pairs and third brightest image for quads.} \label{fig:om10predict} \end{figure} The chance of a given quasar being lensed was determined by OM10 to be $\sim$10$^{-3.5}$, which is comparable to what was obtained with SQLS, that shows a rate of quasar lensing of $\sim$10$^{-3.3}$. In Fig.~\ref{fig:om10predict} we show the expected number of lenses in the full DES survey area as a function of $i$ band magnitude according to the predictions by OM10. We show the expected numbers for pairs and quads (lenses with four quasar images) according to the magnitude of the brightest image in the system ($i_{1}$). For comparison, we also plotted the expected number of lenses according to the magnitude of the fainter image (for pairs) or the third brightest image (for quads), $i_{2/3}$, which is the limit used by OM10. With 50 lensed quasar systems with $i_{1}<19.0$ expected in DES, the challenge becomes how to identify them. SQLS started from a spectroscopic sample of quasars. However, the relatively low numbers of confirmed quasars in the Southern Hemisphere sky and the fact that a large spectroscopic survey is not planned for the next few years, requires a method to photometrically select quasars to look for lenses. Traditionally, that selection would rely on the use of the $u$ band to look for UVX objects \citep[e.g.][]{Croom+2001, Richards+2002}. The lack of this band in DES requires the use of the near and mid IR to make the selection. The use of mid-IR has been applied efficiently for flux limited quasar selection \citep[e.g.][]{Stern+2012, Assef+2013} and \cite{DiPompeo+2015} has shown that the use of SDSS+WISE photometry provided results similar to those obtained with SDSS+UV+near-IR data. % Here, we present results of a search for gravitationally lensed quasars from DES Year 1 observations \citep{Diehl+2014}, obtained between 31 August 2013 and 10 February 2014, combined with JK near infra-red observations from the VISTA Hemisphere Survey (VHS - \citealt{VISTA}; ESO Observing Programme 179.A-2010) and Wide Infra-red Survey Explorer (WISE - \citealt{WISE}). All magnitudes are quoted on the AB system. The conversions from Vega to AB that have been used for the VISTA data are: $J_{AB} = J_{Vega} + 0.937$ and $Ks_{AB} = Ks_{Vega} + 1.839$. These are taken from the Cambridge Astronomical Survey Unit’s website\footnote{\url{http://casu.ast.cam.ac.uk/surveys-projects/vista/technical/filter-set}}. The conversions for the ALLWISE data are $W1_{AB} = W1_{Vega}+2.699$ and $W2_{AB} = W2_{Vega}+3.339$ which are given in \citet{Jarrett+2011} and in the \textit{ALLWISE} explanatory supplement\footnote{The \textit{ALLWISE} explanatory supplement, \url{http://wise2.ipac.caltech.edu/docs/release/allwise/expsup/sec5\_3e.html}, directs the reader to the \textit{WISE All-Sky} explanatory supplement for the conversions; \url{http://wise2.ipac.caltech.edu/docs/release/allsky/expsup/sec4\_4h.html\#summary}.}. When required, a flat cosmology with $\Omega_{m} = 0.3$ and $H_{0} = 70.0$kms$^{-1}$Mpc$^{-1}$ was used unless otherwise specified.
We have identified a high redshift lensed quasar, VDES J2325-5229, by applying GMM supervised machine learning to select the candidates in a colour space defined by DES+VHS+WISE photometric bands. Since the selection does not depend on the $u$-band, we are capable of selecting candidates at higher redshifts. For comparison, amongst the 49 new lenses found by SQLS, only 6 have source redshifts greater than J2325, 3 of which were serendipitously discovered \citep{Johnston+2003,Pindor+2004,McGreer+2010}. Two more lenses with source redshifts greater than J2325 were discovered in SDSS-III BOSS quasar lens survey \citep{More+2016}, including the highest redshift multiply lensed quasar known, with $z_{s}=4.819$. Given the geometry of the system, with the presence of an obvious LRG galaxy, and a angular separation between the quasars that converts to a physical distance of $\sim$23.2kpc, it is unlikely that we are looking at distinct binary quasars. The quasar SEDs are considerably similar, with emission line redshifts that differ by $562\pm240$km/s in the quasar rest frame. Given the poor seeing during the NTT spectroscopic observations and the contamination from the lensing galaxy particularly in the SED of quasar component A, that discrepancy is not surprising. Both quasar images present an intrinsic absorption line system detected in CIV and NV, blue shifted by $2739\pm254$km/s in component A and $2744\pm324$km/s in component B with respect to the emission line redshift, further evidencing the SED similarity. There is evidence that the CIV absorption line is weaker in image B, which would be consistent with different sight lines in the broad-line region of the source quasar as predicted by \cite{Misawa+2014}. This can be seen in Fig~\ref{fig:ntttspectra}, where we have scaled the flux of image B by the median flux ratio between the two quasars blueward of $4000\AA$, where the contamination of the lensing galaxy is most negligible, and subtracted B from A. A strong absorption feature remains at $\lambda\sim5734\AA$, where the CIV absorption is expected to be for $z_{ab}=2.7$. A lesser effect is seen at $\lambda\sim4596\AA$, where one expects the NV absorption to be. This is unsurprising giving the ionization potentials of 64.5eV and 97.9eV for CIV and NV respectively, which means that NV absorbers would be closer to the flux source, and therefore be less likely to be affected by differences in sight lines. Further high-resolution spectroscopy of J2325 would allow for the different sight lines scenario to be confirmed and to constrain the size of the absorber and geometry of the broad-emission line region. The modelled $i$-band magnitude of the lensing galaxy can be used to estimate the rest frame $R$ band magnitude. At these redshifts, and given our choice of rest and observed frame filters, the $K$-corrections can be neglected. For the LRG in J2325-5229, we calculate $M_{R}=-22.41$, given $z=0.4$ and $m_{i}=18.50$. We can compare this magnitude to that obtained by using the separation of the quasar images and the lensing Faber-Jackson relation with fit parameters described in \cite{Rusin+2003} following what was done in \cite{Jackson+2008}. Assuming $z=0.400\pm0.002$ and $\theta=2.90$, \begin{equation} M_{R}=M_{\star R}+2.5\gamma_{E+K}z-1.25\gamma_{FJ}\log{\theta} \end{equation} \noindent yields $M_{R}=-23.11\pm0.53$, which is close to the expected value obtained from modelling the data. The next DES data release will contain data from the full 5000 deg$^{2}$ survey area and will provide a rich sample in which to look for lensed quasars. With such a large area, one can expect to find dozens of bright lenses, including those with quadruple images and the technique introduced in this paper should be able to select all of them as candidates.
16
7
1607.01391
We present the discovery and preliminary characterization of a gravitationally lensed quasar with a source redshift z<SUB>s</SUB> = 2.74 and image separation of 2.9 arcsec lensed by a foreground z<SUB>l</SUB> = 0.40 elliptical galaxy. Since optical observations of gravitationally lensed quasars show the lens system as a superposition of multiple point sources and a foreground lensing galaxy, we have developed a morphology-independent multi-wavelength approach to the photometric selection of lensed quasar candidates based on Gaussian Mixture Models (GMM) supervised machine learning. Using this technique and gi multicolour photometric observations from the Dark Energy Survey (DES), near-IR JK photometry from the VISTA Hemisphere Survey (VHS) and WISE mid-IR photometry, we have identified a candidate system with two catalogue components with I<SUB>AB</SUB> = 18.61 and I<SUB>AB</SUB> = 20.44 comprising an elliptical galaxy and two blue point sources. Spectroscopic follow-up with NTT and the use of an archival AAT spectrum show that the point sources can be identified as a lensed quasar with an emission line redshift of z = 2.739 ± 0.003 and a foreground early-type galaxy with z = 0.400 ± 0.002. We model the system as a single isothermal ellipsoid and find the Einstein radius θ<SUB>E</SUB> ∼ 1.47 arcsec, enclosed mass M<SUB>enc</SUB> ∼ 4 × 10<SUP>11</SUP> M<SUB>⊙</SUB> and a time delay of ∼52 d. The relatively wide separation, month scale time delay duration and high redshift make this an ideal system for constraining the expansion rate beyond a redshift of 1.
false
[ "multiple point sources", "lensed quasar candidates", "high redshift", "SUB", "month scale time delay duration", "a foreground lensing galaxy", "machine learning", "Gaussian Mixture Models", "two blue point sources", "a foreground early-type galaxy", "mass M<SUB", "an elliptical galaxy", "the point sources", "a lensed quasar", "gravitationally lensed quasars", "an emission line redshift", "IR JK", "a foreground z<SUB", "a morphology-independent multi-wavelength approach", "AB</SUB" ]
13.865954
4.331763
79